{ "1002/1002.3656_arXiv.txt": { "abstract": "We made high-resolution spectroscopic observations of limb-spicules in H$\\alpha$ using the Vertical Spectrograph of Domeless Solar Telescope at Hida Observatory. While more than half of the observed spicules have Gaussian line-profiles, some spicules have distinctly asymmetric profiles which can be fitted with two Gaussian components. The faster of these components has radial velocities of 10--40 km~s$^{-1}$ and Doppler-widths of $\\sim$0.4 \\AA \\ which suggest that it is from a single spicule oriented nearly along the line-of-sight. Profiles of the slower components and the single-Gaussian type show very similar characteristics. Their radial velocities are less than 10 km~s$^{-1}$ and the Doppler-widths are 0.6--0.9 \\AA. Non-thermal ``macroturbulent'' velocities of order 30 km~s$^{-1}$ are required to explain these width-values. ", "introduction": "Spicules are one of the fundamental elements of the quiet solar chromosphere, so it is important to know the physical properties of spicules to understand the quiet Sun. \\citet{Pasachoff-Beckers} did a pioneering observation of limb spicules with many chromospheric emission lines, and derived the basic morphological and spectroscopic properties. They also suggested the rotation of spicule gas from the inclined spectra. \\citet{Krat} also made spectroscopic observations of limb spicules in 5 spectral lines, and reported that emission-line profiles are classified into narrow ones and wide ones. However, they were not sure whether they observed single individual spicules, since the spatial resolution of their study was several arcseconds. Summarizing these observational works, \\citet{Beckers} reviewed the morphology and spectroscopic properties of spicules and summarized that a spicule has a diameter of 400--1500 km and mean Doppler velocity of 10 km~s$^{-1}$. \\citet{Kuli-Niko} observed 650 H$\\alpha$ line profiles of 25 spicules, and categorized them into two groups. One group is of relatively small intensity and narrow emission profiles, and the other is of brighter and wider ones. The characteristics of the latter group was thought to be due to the overlapping of unresolved spicules. They suggested that the H$\\alpha$ spicules have temperature of 6000 K and non-thermal velocities of $\\sim$25 km~s$^{-1}$. Follow-up spectroscopic observations of limb spicules under much better seeing conditions were done by \\citet{Kuli} with the 53-cm Lyot coronagraph of Abastumani Astrophysical Observatory and by \\citet{Hasan} with Sacramento Peak VTT. From the study of line-of-sight (LOS) velocity and line-width distribution along the axis of spicules, they got a result that, in majority of spicules, the H$\\alpha$ line-width does not change spatially along the axes of spicules. \\citet{Kuli} suggested that the line width of the majority of spicules are due to the overlapping of gas moving with relative velocities of 20--30 km~s$^{-1}$. \\citet{Nishikawa} used H$\\alpha$ filtergrams to determine the diameter and the motion of limb spicules. This morphological study found that the speed of the apparent up-and-down motion of large spicules is $\\sim$50 km~s$^{-1}$. Assuming ballistic motion, he estimated the initial velocity to be around 100 km~s$^{-1}$, though there has been no other report of such high velocities. The typical diameter of spicules was about 500 km, although the best images show thinner components of 200 and 350 km in diameter. Recent studies using Ca\\emissiontype{II} H filtergram from Hinode has shown that spicules move more dynamically and have much smaller diameters than previously thought. \\citet{DePontieu-b} found that the apparent maximum velocity of ``type I'' spicules is 15--35 km~s$^{-1}$, and the apparent upward velocity of ``Type II'' spicules is 40--300 km~s$^{-1}$. \\citet{Suematsu} reported that a spicule consists of highly dynamic multi-threads as thin as a few tenths of an arcseconds and shows lateral movement or oscillation with rotation. Although Hinode's direct imaging provides us highly valuable images, we need spectroscopic observations of spicules with high spatial and temporal resolutions to know the actual physical state of the spicular gas. As \\citet{Sterling} pointed out, knowledge of the actual physical properties is essential to theoretically clarify the dynamics and ejection process of spicules with numerical simulations. Recently, \\citet{Pasachoff} presented the ground-based observations with very high spatial resolution done at the Swedish 1-m Solar Telescope on La Palma. They analyzed the imaging observations at five wavelengths around H$\\alpha$. Although their observation provides us a new statistics on morphological and dynamical properties of limb spicules, they have derived the spectroscopic quantities such as LOS velocities under the rather simplified assumption of the Gaussian emission profiles. At present, further analyses of full line profiles in detail remain to be done to find the thermodynamical properties of spicular gas. In our series of papers, we present the analyses of spectroscopic observations of limb spicules in H$\\alpha$ with high spatial resolution, which is comparable to the best resolution of ground-based observations. This paper reports and discusses the line widths and the radial velocities of spicules derived from our spectral images. ", "conclusions": "\\label{sec:discuss} We made high-resolution spectroscopic observations of limb spicules in H$\\alpha$. While more than half of the observed spicules have Gaussian line profiles (Type A), some spicules have distinctly asymmetric profiles which can be fitted with two Gaussian components (Type C). The faster component of Type C profiles have radial velocities of 10--40 km~s$^{-1}$ and Doppler width of $\\sim$0.4 \\AA . This width can be explained with temperature $\\simeq$15000 K and microturbulence $\\simeq$5 km~s$^{-1}$. On the other hand, widths of Type A profiles and the slower component of Type C profiles are 0.6--0.9 \\AA. Macroturbulent velocities of order 30 km~s$^{-1}$ are required to explain this large width. Now we will discuss the physical interpretations of these results. Figure \\ref{fig:slit} shows that the slit-line intersects many spicules, and it is natural that the LOS intersects several spicules, too. Though it is difficult to observe a single spicule without superposition of other spicules, Type C profile enables us to separate the emission of single spicule as the fast component of double-Gaussian fitted components. Line width of the fast component can be explained with temperature $\\simeq$15000 K and microturbulence $\\simeq$5 km~s$^{-1}$. Broad width of the Type A profiles suggest that many spicules along the LOS contribute to the observed emission. Figure~\\ref{fig:hist}c shows that their equivalent widths are five times or more larger than those of the fast component. The LOS velocity 20--40 km~s$^{-1}$ of the fast component is consistent with the velocity of the apparent up-and-down motion of spicules reported by many authors (e.g. \\cite{Nishikawa}; \\cite{DePontieu-b}; \\cite{Pasachoff}). Because most spicules are nearly vertical to the solar limb \\citep{Beckers}, their LOS velocities are much smaller than their actual velocities. We suppose that the profiles of Type A and slow component of Type C are from these nearly vertical spicules, and that the fast component is from the few spicules which are almost parallel to the LOS. As shown in section \\ref{sec:width}, we need additional macroturbulent i.e.\\ non-thermal random velocities of order 30 km~s$^{-1}$ to explain the Doppler width of the slow component of Type C and the Type A profiles. Other observations also reported the high non-thermal broadening. For example, \\citet{Makita} suggested that spicules have a turbulence of $\\sim$20 km s$^{-1}$ in the active region, based on the analysis of Ca\\emissiontype{II} H and K profiles in the flash spectrum of the 1958 eclipse. And \\citet{Mariska} reported average non-thermal velocity of 28 km~s$^{-1}$ in EUV line profiles at the quiet limb. Next we proceed to discuss the source of macroturbulence. One of the possible sources is ejecting motion of spicules inclined toward or away from the observer. Variations in speed and inclination of ejection produce the dispersion of LOS velocities. As discussed above, most spicules have small inclinations and so their LOS velocities will be much smaller than 10--40 km s$^{-1}$. Distributions of apparent inclinations reported by \\citet{Pasachoff} have peaks at 10 and 25 degrees, in which case the LOS velocity of ejection speed 40 km s$^{-1}$ will be at most 7 and 17 km s$^{-1}$. They are too small to explain the observed widths. There is also the height difference of ejection speed. The spicules nearer to us or farther than the limb of the sun would be observed in superposition against lower heights of spicules whose bases are on the limb. However, the speed at the tops is no faster than the lower part (\\cite{Sterling}; \\cite{DePontieu-b}), and it will contribute little to the line broadening, similar to the argument above. Another possible source is the lateral motion of spicules due to the Alfv\\'en wave disturbance recently found in Hinode/SOT Ca\\emissiontype{II} H filtergrams (\\cite{DePontieu-a}; \\cite{Suematsu}; \\cite{He}), or the kink wave observed by \\citet{He2}. \\citet{DePontieu-a} showed that the distribution of transverse displacements of the spicules agrees with the velocity amplitudes around 20 km~s$^{-1}$, while \\citet{He} reported the velocity amplitude of high-frequency Alfv\\'en waves to be 4.7--20.8 km~s$^{-1}$. Velocity amplitude of the kink wave reported by \\citet{He2} is less than 8 km~s$^{-1}$. Superposition effects of these motions can broaden the line profiles. Because spicules are believed to be ejected along magnetic field lines, transverse Alfv\\'en waves broaden the H$\\alpha$ line profiles of Type A or slow component of Type C. MHD simulation also showed that the amplitude of Alfv\\'en waves at the chromospheric height to be $\\sim$20 km s$^{-1}$ (\\cite{Suzuki}). However, when the sinusoidal motions of independently disturbed spicules are superposed, its standard deviation will be $1/\\sqrt{2} $ of amplitude. Thus oscillation of velocity amplitude 20 km s$^{-1}$ will contribute only 14 km~s$^{-1}$ to the macroturbulent velocity. As for the fast components, we assume that their surrounding magnetic field is nearly parallel to the LOS, so the Alfv\\'en waves will not contribute to line-broadening, agreeing with their narrow widths even if they are composed of multiple finer threads. \\citet{Kuli} reported the oscillations of the radial velocities of limb spicules with the period of 3--7 min. However, their amplitudes were less than 10 km s$^{-1}$, which are too small to contribute to our observed line-widths. Even when we include superposition effects of ejecting motions and Alfv\\'en wave broadening, it is still insufficient to explain the macroturbulent velocities of $\\sim$30 km~s$^{-1}$. So there must be some additional but unidentified sources for the non-thermal broadening of spicule emission profiles. Alfv\\'en waves with higher frequencies and shorter wavelengths than those observed until now may be present and contribute to the broadening. Let us mention two interesting events found in our analysis. Their characteristics are different from the common spicules discussed so far. While most of radial velocities are less than 40 km~s$^{-1}$, these two cases have much larger ( $>$ 45 km~s$^{-1}$) radial velocities as seen in the upper right area of figure~\\ref{fig:corr}. Their Doppler widths are exceptionally large, too. They are in the spectra at 06:16:44 UT and 06:23:34 UT, and their faint and broad profiles can be found in figure~\\ref{fig:typeC}. The positions of these events are marked in figure \\ref{fig:slit} where there is a gap in the bush of spicules, and where the slit is the nearest to the limb. We need different broadening mechanism for these events. It may be an event with a magnetic reconnection, or something related to the spicule formation. This phenomenon may correspond to the Type II spicules of \\citet{DePontieu-b} or to macrospicules (\\cite{Bohlin}; \\cite{Yamauchi}) or to polar-jets (\\cite{Shibata}; \\cite{Shimojo}; \\cite{Savcheva}). The relations are yet to be ascertained in future studies. In this paper, we treated the results statistically without regard to their positions. However, figure \\ref{fig:spectra} shows that some spicules appear as tilted streaks. This is possibly an evidence of rotation of spicules suggested by \\citet{Pasachoff-Beckers}, and its spatially blurred emission may contribulte to broaden the line profiles. Some physical values are not determined by the H$\\alpha$ profile alone. In order to study the true origin of macroturbulence or of the high-velocity events, we need simultaneous spectroscopic observations of spicules in multiple spectral lines, which can be achieved with the Horizontal Spectrograph at Hida Observatory. We also need high spatial-resolution images of Hinode/SOT simultaneously obtained with the spectroscopic observation to understand the details of spicules. We will plan our next observation to accommodate to those requirements. \\bigskip We thank the referees for the greatly helpful comments. We are grateful for the use of EIT data obtained on the SOHO spacecraft. SOHO is a project of international cooperation between ESA and NASA. The authors are supported by a grant-in-aid for the Global COE program ``The Next Generation of Physics, Spun from Universality and Emergence'' from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan, and by the grant-in-aid for ``Creative Scientific Research The Basic Study of Space Weather Prediction'' (17GS0208, PI: K. Shibata) from the Ministry of Education, Culture, Sports, Science and Technology of Japan, and also partly supported by the grant-in-aid from the Ministry of Education, Culture, Sports, Science and Technology of Japan (No.19540474)." }, "1002/1002.1715_arXiv.txt": { "abstract": "I describe new constraints on the lifetimes and morphologies of transitional protoplanetary disks from observations of 1--10 Myr old stars with the \\textit{Spitzer Space Telescope}. New Spitzer results clearly show evidence for two kinds of transitional disks and thus two main disk evolutionary pathways: disks which form an inner hole/gap and clear from the inside out and disks that deplete more homologously. Analyzing the disk populations of 1--10 Myr old clusters such as Taurus, IC 348, NGC 2362, and $\\eta$ Cha show that the mean transitional disk lifetime must be an appreciable fraction of the mean protoplanetary disk lifetime: $\\approx$ 1 Myr out of 3--5 Myr. The varieties of transitional disk SEDs and correlations with other disk diagnostics are consistent with multiple mechanisms responsible for clearing disks. ", "introduction": "\\textit{Transitional} protoplanetary disks bridge the evolutionary gap between luminous optically-thick \\textit{primordial} disks of gas and small dust, which presumably have yet to make planet-mass bodies and gas-poor/free optically-thin \\textit{debris} disks, which have ended any gas giant planet formation \\citep{Strom1989, CurrieLada2009}. Stars surrounded by transitional disks have near-to-mid IR dust emission intermediate between primordial disk-bearing stars and diskless photospheres, implying that much of their solid mass is in the process of being lost from the system and/or incorporated into large planetesimals/protoplanets. Ground-based studies of transitional disks from IR to submm photometry find evidence for structural features in these disks, large inner holes/gaps in the disks' dust distribution, indicative of active disk dispersal \\citep[e.g.][]{Calvet2002}. Thus, transitional disks may provide valuable insights into how and when planet formation ends. The unprecedented mid-IR sensitivity of the \\textit{Spitzer Space Telescope} made spectroscopic observations of nearby transitional disks and photometry for many transitional disks beyond $\\sim$ 200--400 pc accessible for the first time. In this contribution, I summarize new Spitzer results that clarify our understanding of morphologies and lifetimes of transitional disks. These results reveal two types of transitional disks, which may be evidence for a range of mechanisms responsible for dispersing disks and show that the transitional disk phase typically comprises an appreciable fraction of the total protoplanetary disk lifetime. Furthermore, analyzing the accretion frequency/rate and submillimeter fluxes of transitional disks provides some insight into the processes that are plausibly responsible for transitional disk morphologies and thus potentially crucial for dispersing disks and shutting off planet formation. ", "conclusions": "" }, "1002/1002.0525_arXiv.txt": { "abstract": "Assuming the standard cosmological model as correct, the average linear size of galaxies with the same luminosity is six times smaller at $z=3.2$ than at $z=0$, and their average angular size for a given luminosity is approximately proportional to $z^{-1}$. Neither the hypothesis that galaxies which formed earlier have much higher densities nor their luminosity evolution, mergers ratio, or massive outflows due to a quasar feedback mechanism are enough to justify such a strong size evolution. Also, at high redshift, the intrinsic ultraviolet surface brightness would be prohibitively high with this evolution, and the velocity dispersion much higher than observed. We explore here another possibility to overcome this problem by considering different cosmological scenarios that might make the observed angular sizes compatible with a weaker evolution. One of the models explored, a very simple phenomenological extrapolation of the linear Hubble law in a Euclidean static universe, fits the angular size vs. redshift dependence quite well, which is also approximately proportional to $z^{-1}$ with this cosmological model. There are no free parameters derived ad hoc, although the error bars allow a slight size/luminosity evolution. The type Ia supernovae Hubble diagram can also be explained in terms of this model with no ad hoc fitted parameter. WARNING: I do not argue here that the true Universe is static. My intention is just to discuss which theoretical models provide a better fit to the data of observational cosmology. ", "introduction": "The analysis of the dependence of the angular size of some sources with redshift was for many decades one of the most important geometric tests of cosmological models. Different cosmologies predict different dependences for a given linear size and this can be compared with the data from observations. The test, first conceived by Hoyle\\cite{Hoy59}, is simple in principle but its application is not so simple because of the difficulty in finding a standard rod, a type of object with no evolution in linear size over the lifetime of the Universe. It is well known that the application of the angular size ($\\theta $) vs. redshift ($z$) test gives a rough dependence of $\\theta \\propto z^{-1}$ for QSOs and radio galaxies at radio wavelengths\\cite{Mil71,Kel72,War74,Kap77,Kap87,Uba93}, for first ranked cluster galaxies in the optical\\cite{San72,Djo81,Pas96,Sch97}, and for the separation of brightest galaxies in clusters \\cite{LaV86} or in QSO-galaxy pairs of the same redshift\\cite{Sap99}. The deficit of large objects at high redshifts with respect to the predictions of an expanding Universe is believed to be an evolutionary effect by which galaxies were smaller in the past (e.g., \\cite{Mil71}), or a selection effect (e.g., \\cite{Jac73}). Thus, the $\\theta \\propto z^{-1}$ relationship, as predicted by a static Euclidean Universe, would be just a fortuitous coincidence of the superposition of the $\\theta (z)$ dependence in the expanding Universe and evolutionary/selection effects. Some other studies have tried to find better standard rods. Ultra-compact radio sources\\cite{Kel93,Jac97,Gur99,Jac04,Jac06} were used to carry out an angular size test: a dependence different from $\\theta \\propto z^{-1}$ and closer to the predictions of expanding Universe models was found. The test was even used to ascertain not only whether or not the Universe is expanding but also to constrain the different cosmological parameters. However, these applications are not free from selection effects\\cite{Jac04} and, as will be discussed in \\S \\ref{.ultracom}, interpretation of the results of these tests is not so straightforward. Another proposal\\cite{Mar08a} used the rotation speed of high redshift galaxies as a standard size indicator since there is a correlation between size and rotational velocity of galactic disks. This method was indeed applied\\cite{Mar08b} over a sample of emission-line galaxies with $0.2 4$, and is a member of a rich proto-cluster of galaxies at $z = 4.05$ in GOODS-North. We have observed the CO 1-0 and 2-1 emission with the VLA, the CO 6-5 emission with the PdBI Interferometer, and the 5-4 emission with CARMA. The H$_2$ mass derived from the CO 1-0 emission is $1.3 \\times 10^{11} (\\alpha/0.8)$ M$_\\odot$. High resolution imaging of CO 2-1 shows emission distributed over a large area, appearing as partial ring, or disk, of $\\sim 10$kpc diameter. The integrated CO excitation is higher than found in the inner disk of the Milky Way, but lower than that seen in high redshift quasar host galaxies and low redshift starburst nuclei. The CO 4-3 integrated line strength is more than a factor of two lower than expected for thermal excitation. The excitation can be modeled with two gas components: a diffuse, lower excitation component with a radius $\\sim 4.5$kpc and a filling factor $\\sim 0.5$, and a more compact, higher excitation component (radius $\\sim 2.5$kpc, filling factor $\\sim 0.13$). The lower excitation component contains at least half the molecular gas mass of the system, depending on the relative conversion factor. The VLA CO 2-1 image at $0.2\"$ resolution shows resolved, clumpy structure, with a few brighter clumps with intrinsic sizes $\\sim 2$ kpc. The velocity field determined from the CO 6-5 emission is consistent with a rotating disk with a rotation velocity of $\\sim 570$ km s$^{-1}$ (using an inclination angle of 45$^o$), from which we derive a dynamical mass of $3 \\times 10^{11}$ \\msun within about 4 kpc radius. The star formation distribution, as derived from imaging of the radio synchrotron and dust continuum, is on a similar scale as the molecular gas distribution. The molecular gas and star formation are offset by $\\sim 1\"$ from the HST I-band emission, implying that the regions of most intense star formation are highly dust-obscured on a scale of $\\sim 10$kpc. The large spatial extent and ordered rotation of this object suggests that this is not a major merger, but rather a clumpy disk accreting gas rapidly in minor mergers or smoothly from the proto-intracluster medium. Qualitatively, the kinematic and structural properties of GN20 compare well to the most rapid star-formers fed primarily by cold accretion in cosmological hydrodynamic simulations. Conversely, if GN20 is a major, gas rich merger, then some process has managed to ensure that the star formation and molecular gas distribution has not been focused into one or two compact regions. ", "introduction": "Studies of the stellar populations of elliptical galaxies imply that massive ellipticals form the bulk of their stars fairly quickly (timescales $\\le 1$ Gyr) at early epochs ($z > 2$). Moreover, there is a clear trend with increasing mass such that the more massive the galaxy, the earlier and quicker the star formation (see review by Renzini 2006). This conclusion is supported by studies of specific star formation rates (SFR/stellar mass), indicating `downsizing' in galaxy formation, with active star formation being preferentially quenched in more massive galaxies over cosmic time (Noeske et al. 2007, 2009; Zheng et al. 2007; Pannella et al. 2009), as well as the direct observation of old stellar populations in early type galaxies at $z \\ge 1$, implying formation redshifts $z > 3$ (Collins et al. 2009; Kurk et al. 2009, Kotilainen et al. 2009; Papovich et al. 2010). These results imply that there should be a progenitor population of active, clustered star forming galaxies at high redshift. Bright submm-selected galaxies (SMGs; $S_{850\\mu m} > 5$ mJy; see Blain et al. 2002 for a review) are an important class of source in this regard. Although they are relatively rare, with typical space densities of $10^{-5}$--$10^{-6}$~Mpc$^{-3}$, they have very high bolometric luminosities ($\\sim 10^{13}$ \\lsun), implying the most intense bursts of star formation known ($\\sim 1000$ \\msun yr$^{-1}$). While radio galaxies and quasars can reach similar or higher luminosities (Miley \\& de Breuck 2007; Solomon \\& vanden Bout 2006), the latter objects contain bright active galactic nuclei (AGN), and it is possible that some of their far-infrared emission is powered by accretion onto black holes. SMGs, on the other hand, are known to be largely star formation dominated galaxies (Alexander et al. 2005). The emerging scenario is that SMGs may be the starburst progenitors of massive early type galaxies. These hyper-luminous high-z galaxies often trace high overdensities (Stevens et al. 2003; Aravena et al. 2009; although cf. Chapman et al. 2009), and are likely related to the formation of giant elliptical galaxies in clusters. A key question for the SMGs is: what drives the prolific star formation? Tacconi et al. (2006; 2008) argue, based on CO imaging of a sample of $z \\sim 2$ SMGs, that SMGs are predominantly nuclear starbursts, with median sizes $< 0.5\"$ ($< 4$kpc), 'representing extreme, short-lived, maximum star forming events in highly dissipative mergers of gas rich galaxies.' This conclusion is supported by VLBI imaging of the star forming regions in two SMGs (Momjian et al. 2005; 2009; 2010). We return to this question below. The discovery of apparently old elliptical galaxies at $z \\ge 2$ has pushed the question of starburst progenitors of giant elliptical galaxies to even earlier epochs (Cimatti et al 2004; Wikind et al. 2008; Mobasher et al. 2005; Kriek et al. 2008; Doherty et al. 2009). The redshift distribution for about 50\\% of the SMG population, namely, the radio detected sources, has been shown to peak around $z \\sim 2.3$, with most of these sources being between $z \\sim 1.5$ and 3 (Carilli \\& Yun 2000; Chapman et al. 2003; Wagg et al. 2009). However, there is a low redshift bias in radio-selected samples, and the question remains: is there a substantial ($\\sim 30\\%$) population of SMGs at $z > 3$? These higher redshift sources would potentially pin-point very early formation of the most massive ellipticals. Early searches for SMGs at $z>4$ (Dannerbauer et al. 2002; 2004; 2008; Dunlop et al. 2004; Younger et al. 2007; 2008; Wang et al. 2007), were unsuccessful. However, recently, a number of SMGs have been found at $z > 4$, including two in the COSMOS field ($z = 4.5$ and 4.7; Capak et al. 2008; Schinnerer et al. 2008, 2009), GN10 at $z \\sim 4.04$ in GOODS-North (Daddi et al. 2009b), a $z = 4.76$ SMG in the CDF-South (Coppin et al. 2009), a strongly lensed source at $z = 4.044$ (Knudsen et al. 2009), and GN20, the subject of this paper. Daddi et al. (2009a) conclude, based on SMG space densities and duty cycles, that there are likely enough SMGs at $z > 3.5$ to account for the known populations of old massive galaxies at $z \\sim 2$ to 3. They also point out that the contribution of SMGs to the comoving cosmic star formation rate density at $z \\sim 4$ (SFRD $\\sim 0.02$ M$_\\odot$ yr$^{}$ Mpc$^{-3}$) is comparable to that of Lyman break galaxies. ", "conclusions": "\\subsection{GN20} We have presented the most detailed imaging analysis to date of the CO emission from an SMG, including imaging the lower order transitions down to 1 kpc resolution. The principle physical parameters resulting from this study are listed in Table 1. The main result from this work is that the molecular gas and star formation are well resolved on a scale $\\sim 10$kpc. The high resolution CO 2-1 imaging, in particular, shows a partial ring, or disk, on this scale. The ring shows a few resolved clumps with (deconvolved) sizes $\\sim 2$ kpc, but no single clump dominates the total emission. This is also true for the 1.4 GHz continuum emission, presumably tracing star formation, which has a similar morphology to the CO emission. The higher order CO observations indicate a regular velocity field, consistent with a disk with a rotational velocity of 570 km s$^{-1}$. The total H$_2$ mass derived from the CO 1-0 emission is $1.3\\times 10^{11} \\times (\\alpha/0.8)\\msun$, which is roughly 40\\% of the dynamical mass within 4 kpc radius. The entire $\\sim 10$ kpc region of active star formation, as traced by the CO, FIR, and radio continuum, is completely obscured in the HST I-band (rest-frame UV) image. The CO is lower excitation than seen in low redshift nuclear starbursts and high redshift quasar host galaxies, but it is higher than in nearby spiral galaxies and normal star forming galaxies at $z \\sim 1.5$. The CO emission from GN20 is consistent with a two component model, consisting of a 4.5 kpc radius disk of lower density (300 cm$^{-3}$), temperature (30K), with a filling factor $\\sim 0.5$, and a region of $\\sim 2.5$ kpc radius with higher density ($\\sim 6300$ cm$^{-3}$), higher temperature (45K), and lower filling factor ($\\sim 0.13$). The mass is roughly equal in each component (assuming the same conversion factor). The gas depletion timescale is comparable to the rotational time of the galaxy $\\sim 5 \\times 10^7 \\times (\\alpha/0.8$) years. We note that Papadopoulos et al. (2010) have proposed that dust opacity in dense regions can also affect the observed line ratios for CO, when observing very high order transitions (eg. CO6-5). High resolution imaging of the CO6-5 is required to determine the spatial dependence of gas excitation in GN20. \\subsection{Star formation in GN20} Two mechanisms have been proposed in recent years for driving active star formation in high redshift galaxies: major gas rich mergers (Narayanan et al. 2009) and cold mode accretion (Dekel et al. 2009; Keres et al. 2009; Keres et al. 2005). The process of fueling nuclear starbursts via major gas rich mergers is well studied in the nearby Universe (eg. Mihos \\& Hernquist 1996; Barnes \\& Hernquist 1991). The general idea is that gravitational torques induce strong dissipation and inflow of gas, leading to an increase in the star formation rates by up to two orders of magnitude over quiescent disks on the short timescale of the merger $\\sim$ few$\\times 10^7$ years. Most of this star formation occurs on scales $< 1$ kpc in the galaxy nuclei, as is seen in nearby ULIRGs (Downes \\& Solomon 1998). Major gas rich mergers have been invoked to explain the compact, maximal starbursts seen in some high redshift quasar host galaxies (Li et al. 2007; Walter et al. 2009; Riechers et al. 2009). Johansson et al. (2009) point out, if a nuclear starburst is merger-driven, the mass ratio has to be close to unity. Tacconi et al. (2006; 2008) conclude, based on high resolution CO imaging, that $z \\sim 2$ SMGs are nuclear starbursts on scales $<4$kpc, driven by major gas rich mergers. However, they base this conclusion on observations of high order CO lines (3-2 and higher at sub-arcsecond resolution). In GN20 we see a more extended, lower excitation molecular gas distribution on a scale $\\sim 10$ kpc, containing at least half the gas mass in the system. Interestingly, a number of other $z \\sim 2$ SMGs have been observed in CO 1-0: the submm-bright ERO J16450+4626 (Greve et al. 2003), SMM J13120+4242 (Hainline et al. 2006), and SMM J02399-0136 (Ivison et al. 2010). These galaxies show excess CO 1-0 emission relative to what is expected by extrapolating from higher-order transitions assuming constant brightness temperature, by factors of at least two. Moreover, VLA imaging of J16450+4626 reveals extended CO 1-0 emission on a scale of $\\sim 10$ kpc (Greve et al. 2003), while for J02399-0136 the CO 1-0 emission extends over 25 kpc (Ivison et al. 2010). An alternative model, known as cold mode accretion (CMA), or stream fed galaxy formation, has recently been proposed to explain secular star formation (ie. on timescales $> 10^8$ years) in more populous, normal star forming galaxies at $z \\sim 2$ (Dekel et al. 2009; Keres et al. 2009). In the CMA model, gas flows into galaxies from the IGM along cool, dense filaments. The flow never shock-heats due to the rapid cooling time, but continuously streams onto the galaxy at close to the free-fall time. This gas forms a thick, turbulent, rotating disk which efficiently forms stars across the disk, punctuated by giant clouds of enhanced star formation on scales $\\sim$ few kpc. These star forming regions then migrate to the galaxy center via dynamical friction and viscosity, forming compact stellar bulges (Genzel et al. 2006; Genzel et al. 2008; Bournaud et al. 2008a,b; Elmegreen et al. 2009). The CMA process can lead to relatively steady and active ($\\sim 100$ M$_\\odot$ yr$^{-1}$) star formation in galaxies over timescales approaching 1 Gyr. The process slows down dramatically as gas supply decreases, and the halo mass increases, generating a virial shock in the accreting gas. Subsequent dry mergers at lower redshift then lead to continued total mass build up, and morphological evolution, but little subsequent star formation (Hopkins et al. 2009; Naab et al. 2009). Observations of intermediate redshift ($z \\sim 2$), normal star forming galaxies support the CMA model (Genzel et al. 2006, 2008; Daddi et al. 2008; 2009b; 2010a; Tacconi et al. 2010). Dav\\'e et al (2009) suggest that a substantial fraction of SMGs could be fed primarily by CMA, and not major mergers. Hydrodyanamic simulations show that quite high gas accretion rates can be achieved in large halos at early epochs, and elevations of SFR over the average accretion rate by a factor 2 to 3 can be frequent owing to (common) minor mergers. While the Dav\\'e et al. study focused on $z=2$, Finlator et al. (2006) studied galaxies in cosmological hydrodynamic simulations at $z=4$ and found two instances (within a $3\\times 10^6$ comoving Mpc$^3$ volume) of galaxies forming stars at $>1200$ M$_\\odot$ yr$^{-1}$. These galaxies had stellar masses several~$\\times 10^{11}M_\\odot$, very similar to GN20. They were found to be forming stars at $\\sim 2$ to $2.5\\times$ their average rate, owing to their location at the center of the largest potential wells with constantly infalling satellite galaxies, ie. environments comparable to GN20. While their star formation rates are still a factor $\\sim 2$ to 3 lower than GN20, given the uncertainties in conversion from $L_{IR}$ to SFR (eg. the IMF), and uncertainties in the models, there is at least a plausible association of these simulated galaxies with GN20. The CMA and major merger scenarios might be expected to have different structural and kinematic signatures. The ordered rotation and extended gas distribution would favor a disk that is not being strongly disturbed. The gas distribution (Figure 2) and velocity field (Figure 5) of GN20 are qualitatively similar to that seen in the simulated SMG maps in Figure~5 of Dav\\'e et al. (2009), particularly object ``B\" which is a quiescently star-forming thick disk, though we note that the SFR of this $z=2$ simulated galaxy is substantially lower than GN20. On the other hand, such signatures do not conclusively rule out a merger, as Robertson \\& Bullock (2008) have shown that ordered rotation of an extended gas disk can be reestablished very shortly after a major merger. Therefore we cannot make firm conclusions about the driver of star formation in GN20, although the simplest interpretation of the molecular gas data favors CMA. A key point is that, if GN20 is an on-going major gas rich merger, then some process has managed to ensure that the star formation and molecular gas distribution has not been focused into one or two compact nuclear regions. This latter point also begs the question of the progeny of a system such as GN20, since massive (stellar masses $\\sim 10^{11}$ M$_\\odot$) 'red and dead' ellipticals at $z \\sim 2$ typically show fairly compact stellar distributions, with radii $\\sim 1$kpc (van Dokkum et al. 2008), although questions have been raised about morphological K-corrections and the presense of an AGN (Daddi et al. 2005). On the other hand, star forming galaxies (SFR $\\sim 100$ M$_\\odot$ year$^{-1}$) of this stellar mass at $z \\sim 2$ are often seen to have spatially extended stellar distributions, comparable to the CO size of GN20 (Kriek et al. 2009; Daddi et al. 2008; 2010a). Hence, if GN20 is to evolve into a passive elliptical at $z \\sim 2$, the stellar distribution will have to evolve to a more compact configuration. Conversely, GN20 could remain a star forming galaxy for a long period, although at a substantially lower star formation rate. A number of key observations are required to untangle the mechanisms driving star formation in GN20. First, high resolution imaging of the millimeter continuum will reveal the distribution of star formation across the disk in greater detail. Second, high resolution imaging of the high order transitions can be used to determine the spatial excitation of the molecular gas. And third, high spectral and spatial resolution imaging of the low order CO emission is required to determine the gas dynamics on both large and small scales, ie. verify the overall rotation, and determine the internal turbulent velocity and stability parameters in the disk. Shapiro et al. (2009) show that such a study of disk kinemetry enables an 'empirical differentiation between merging and non-merging systems'. These latter observations have now become possible with the Expanded Very Large Array. \\subsection{Lensing} An open issue remains gravitational lensing, in particular given the partial ring-like morphology of GN20 in CO emission. An Einstein ring has been observed in CO in the strongly lensed, $z = 4.12$ quasar host galaxy J2322+1944 (Carilli et al. 2003; Riechers et al. 2008). Lensing would also help explain the extreme (apparent) luminosity of GN20 (Pope et al. 2006). To date, there is no evidence for a lensing galaxy in the HST image of GN20. This is unusual, given the depth of the HST data, and the $\\sim 1''$ diameter of the ring. For example, the CO Einstein ring in J2322+1944 has a similar diameter as GN20, and the lensing galaxy is seen clearly in the F814W filter image with a magnitude of 21.9. This galaxy would be easily detected in the HST I-band image of GN20. Also, the regular velocity field in the 6-5 line argues against lensing, since caustic structures can lead to complex apparent velocity structures in the image plane, as is the case for J2322+1944 (Riechers et al. 2008). Lastly, the fact that the gas mass and the dynamical mass are comparable (within a factor 2 or so), argues against very strong lensing." }, "1002/1002.1245_arXiv.txt": { "abstract": "{ The dynamics of a barred galaxy depends on the angular velocity or pattern speed of its bar. Indeed, it is related to the location of corotation where gravitational and centrifugal forces cancel out in the rest frame of the bar. The only direct method for measuring the bar pattern speed is the Tremaine-Weinberg technique. This method is best suited to the analysis of the distribution and kinematics of the stellar component in absence of significant star formation and patchy dust obscuration. Therefore, it has been mostly used for early-type barred galaxies. The main sources of uncertainties on the directly-measured bar pattern speeds are discussed. There are attempts to overcome the selection bias of the current sample of direct measurements by extending the application of the Tremaine-Weinberg method to the gaseous component. Furthermore, there is a variety of indirect methods which are based on the analysis of the gas distribution and kinematics. They have been largely used to measure the bar pattern speed in late-type barred galaxies. Nearly all the bars measured with direct and indirect methods end close to their corotation radius, i.e., they are as rapidly rotating as they can be. ", "introduction": "Strong bars are observed in optical images of roughly half of all the nearby disk galaxies \\citep{Marinova2007, Barazza2008, Aguerri2009}. Therefore bars are a common feature in the central regions of disk galaxies. Their growth is partly regulated by the exchange of angular momentum with the stellar disk and the dark matter halo. For this reason the dynamical evolution of bars can be used to constrain the content and distribution of dark matter in the inner regions of galaxy disks \\citep[e.g.,][]{Weinberg1985, Bertin1989, Debattista2000}. The morphology and dynamics of a barred galaxy depend on the angular velocity or pattern speed of the bar, $\\om$. Usually, the bar pattern speed is parametrized with the bar rotation rate $\\vpd\\equiv\\lag/\\len$. This is the distance-independent ratio between the corotation radius $\\lag$, where the gravitational and centrifugal forces cancel out in the rest frame of the bar, and the length of bar semi-major axis $\\len$. The corotation radius is derived from the bar pattern speed as $\\lag = V_{\\rm c}/\\om$, where $V_{\\rm c}$ is the disk circular velocity. As far as the value of $\\vpd$ is concerned, if $\\vpd < 1.0$ the stellar orbits are elongated perpendicular to the bar and the bar dissolves. For this reason, self-consistent bars cannot exist in this regime. Bars with $\\vpd \\gtrsim 1.0$ are close to rotating as fast they can, and there is no a priori reason for $\\vpd$ to be significantly larger than 1.0. The value of $\\vpd$ can be used to classify bars into fast ($1.0 \\leq \\vpd \\leq 1.4$) versus slow ($\\vpd > 1.4$), with the dividing value at 1.4 by consensus \\citep{Debattista2000}. Note the value of $\\vpd$ does not imply a specify value of the pattern speed. ", "conclusions": "All the bars in Table~1 are consistent with being fast (Fig.~\\ref{fig:r}). This is particularly true when galaxies with small uncertainties on $\\vpd$ are considered. In fact, the probability that the bar length is twice as long as the corotation radius ($\\vpd>2$) is $12\\%$ for all the sample galaxies and $8\\%$ for galaxies with $\\Delta \\vpd / \\vpd\\leq0.3$. The median rotation rate is $\\vpd=1.2$. The fact that some of the values of bar rotation rate are nominally $\\vpd<1$ has been interpreted by \\citet{Debattista2003} as due to the scatter introduced by uncertainties on PA$_{\\rm disk}$. The quoted uncertainties on $\\vpd$ are heterogeneous and include both $68\\%$ confidence intervals and maximal errors. It would be very useful if they were calculated and given in an homogeneous way. For example, they could be estimated from Monte Carlo simulations based on the uncertainties on $\\len$ and $\\lag$. No trend in $\\vpd$ is observed with morphological type. But, the sample of bars studied so far with the stellar-based TW method is biased toward the bright and strongly barred SB0 and SB0/a galaxies. One of them hosts two nested bars (NGC~2950, \\citealt{Corsini2003}). The sample includes only 3 spiral galaxies (NGC~271, \\citealt{Gerssen2003}; NGC~2523, \\citealt{Treuthardt2007}; NGC~3992, \\citealt{Gerssen2003}), and one dwarf galaxy (NGC~4431, \\citealt{Corsini2007}). The bar rotation rate of NGC~4431 is $\\vpd = 0.6^{+1.2}_{-0.4}$ at $99\\%$ confidence level \\citep{Corsini2007}. Albeit with large uncertainty, the probability that the bar ends close to its corotation radius is about twice as likely as that the bar is much shorter than the corotation radius. This suggests a common formation mechanism of the bar in both bright and dwarf galaxies. If their disks were previously stabilized by massive dark matter halos, bars were not produced by tidal interactions because they would be slowly rotating \\citep{Noguchi1999}. But, this is not the case even in the two sample galaxies which show signs of weak tidal interaction with a close companion (i.e., NGC~1023, \\citealt{Debattista2002}; NGC~4431, \\citealt{Corsini2007}). There is no difference between TW measurements of the stellar component in isolated or mildly interacting barred galaxies. Besides, neither the length nor strength of the bars are found to be correlated with the local density of the galaxy neighborhoods \\citep{Aguerri2009}. The bars of ESO~139-G09 \\citep{Aguerri2003} and NGC~1358 \\citep{Gerssen2003} are weak and fast. Thus, the hypothesis of \\citet{Kormendy1979} that weak bars are the end state of slowed down fast bars is not supported by observations. They instead favor a scenario in which weak and strong bars form in the same way. The $\\vpd$ determinations based on the stellar TW method agree with those obtained by indirect methods, which are largely adopted for gas-rich galaxies. According to the compilations by \\citet{Elmegreen1996} and \\citet{Rautiainen2008}, almost all the measured bars have $1.0 \\leq \\vpd \\leq 1.4$ within the errors. The same is true also for the $\\vpd$ values obtained from pattern speeds measured with the gas-based TW method \\citep{Bureau1999, Fathi2009, Chemin2009}. A fast bar is ruled out by errors only in NGC~2917 \\citep[$\\vpd>1.7$,][]{Bureau1999} and UGC~628 \\citep[$\\vpd=2.0^{+0.5}_{-0.3}$,][]{Chemin2009}. If bars of gas-poor lenticulars and early-type spirals have the same $\\vpd$ as gas-rich late-type spirals, then gas is not dynamically important for the evolution of bar pattern speed \\citep{Debattista2003}. Unfortunately, uncertainties on the measured $\\vpd$ are often not quoted for late-type galaxies. This missing piece of information is crucial to derive the $\\vpd$ distribution as a function of the morphological type. Additional work is needed for the stellar TW method to increase both the accuracy of the $\\om$ measurements and extend the number of studied late-type barred galaxies. Integral-field spectroscopy overcomes many problems of long-slit observations and leads to more efficient and accurate TW measurements. Dwarf barred galaxies are ideal targets to this aim. They nicely fit the field of view of the available integral-field spectrographs and are good candidates for testing the predictions that dark-matter dominated barred galaxies should have a slow bar. A successful application of the TW method to the stellar component of late-type barred galaxies would remedy the selection bias present in the current sample of measured $\\om$. This will allow a straightforward comparison of the results of stellar-based TW method with indirect and gas-based TW measurements in the same range of Hubble types." }, "1002/1002.3076_arXiv.txt": { "abstract": "{Studying outliers from the bimodal distribution of galaxies in the color-mass space, such as morphological early-type galaxies residing in the blue cloud (\\emph{blue E/S0s}), can help to better understand the physical mechanisms that lead galaxy migrations in this space.} { In this paper we try to bring new clues by studying the evolution of the properties of a significant sample of blue E/S0 galaxies in the COSMOS field. } {We define blue E/S0 galaxies as objects having a clear early-type morphology on the HST/ACS images (according to our automated classification scheme \\textsc{galSVM}) but with a blue rest-frame color (defined using the SED best fit template on the COSMOS primary photometric catalogues). Combining these two measurements with spectroscopic redshifts from the zCOSMOS 10k release, we isolate 210 $I_{AB}<22$ blue early-type galaxies with $M_*/M_\\odot>10^{10}$ in three redshift bins ($0.210^{10}$ from $z\\sim1$ in the COSMOS field. Our sample is complete up to $M_*/M_{\\odot}\\sim10^{10}$ in the lower redshift bin and the completeness limit rises to $M_*/M_{\\odot}\\sim5\\times10^{10}$ at $z\\sim1$. Spectroscopic redshifts come from the zCOSMOS 10k release. Morphologies are determined in the ACS images using our tested and validated code \\textsc{galSVM} and the spectral classification is performed using a best-fitting technique with 38 templates on the primary photometric COSMOS catalogues. We define a blue E/S0 galaxy as an object presenting an early-type morphology and which SED best fits with a blue template. \\\\ The main properties of these galaxies summarize as follows: \\begin{itemize} \\item Globally blue E/S0 galaxies represent $\\sim5\\%$ of the whole sample of early-type galaxies with $M_*/M_{\\odot}>10^{10}$ from $z\\sim1.4$. Nevertheless the abundance strongly depends on the considered stellar mass. Below a threshold mass ($M_t$), they become more abundant and represent about $20\\%$ of the red early-type population. This threshold mass evolves with redshift from $M_t\\sim11.1$ at $z\\sim1$ to $M_t\\sim10.1$ at $z\\sim0.2$. This evolution matches the evolution of the bimodality mass (crossover mass between late-type blue population and the early-type red population). Both masses are therefore a measure of the build-up of the RS. In contrast, the shutdown mass (mass above which blue galaxies dissappear) does not evolve. This could be related with differences in the gas physics above a critical mass. \\item Blue E/S0 galaxies fall in the same region in the magnitude-radius space than blue compact galaxies defined by \\cite{guzman97}. However it is not the same population since we focus here in the massive tail of the distribution which is dominated by passive red E/S0 galaxies. \\item The visual inspection of these galaxies does not reveal important perturbations up to out resolution limit. On the contrary, galaxies seem to have a well-defined shape with high-central light concentration. Low mass galaxies seem to have more small companions than the massive counter-parts, however we do not have spectroscopic information to confirm this trend. \\item Massive blue E/S0 ($log{M_*/M_\\odot}>10.5$) galaxies tend to have similar sizes than E/S0 galaxies. At lower masses however they tend to lie between spirals and spheroids. \\item The specific star formation rate measured with the [OII]3727 line and stellar masses, shows that blue E/S0 galaxies have similar SSFR values than normal star-forming spirals. The outliers of this relation are galaxies more massive than $log(M_*/M_\\odot)\\sim10.8$. These outliers present SSFR characteristic of post-merger galaxies ($\\sim 10^{2.5}$ $Gyr^{-1}$). \\end{itemize} These results seem to point out that these galaxies have different natures depending on their stellar mass. Massive objects ($log(M_*/M_\\odot)>10.8$), are probably post mergers galaxies migrating to the red-sequence while less massive galaxies are more likely progenitors of todays late-type galaxies which are accreting gas. The migration of the massive galaxies is faster than $3$ Gyr as expected by numerical simulations of merger remnants. Indeed, the evolution of the number density computed in redshift bins of equal time show that the most massive E/S0 galaxies evolved in typical time-scales smaller than 3 Gyr. In contrast the evolution of blue early-type galaxies with smaller masses is produced in larger time-scales. This low-mass galaxies might be the result of minor-mergers with surrounding satellites which triggered the central star-formation and built-up a bulge component, and therefore appear with an early-type shape. \\\\ The confirmation of these hypothesis requires a detailed analysis of the internal kynematics of these objects to see if the physics is indeed different. In particular, 3D spectroscopy should be useful to get more details about their nature." }, "1002/1002.2136_arXiv.txt": { "abstract": "{Cosmic shear is considered one of the most powerful methods for studying the properties of Dark Energy in the Universe. As a standard method, the two-point correlation functions $\\xi_\\pm(\\vt)$ of the cosmic shear field are used as statistical measures for the shear field.}{In order to separate the observed shear into E- and B-modes, the latter being most likely produced by remaining systematics in the data set and/or intrinsic alignment effects, several statistics have been defined before. Here we aim at a complete E-/B-mode decomposition of the cosmic shear information contained in the $\\xi_\\pm$ on a finite angular interval.}{We construct two sets of such E-/B-mode measures, namely Complete Orthogonal Sets of E-/B-mode Integrals (COSEBIs), characterized by weight functions between the $\\xi_\\pm$ and the COSEBIs which are polynomials in $\\vt$ or polynomials in $\\ln\\vt$, respectively. Considering the likelihood in cosmological parameter space, constructed from the COSEBIs, we study their information content.}{We show that the information grows with the number of COSEBI modes taken into account, and that an asymptotic limit is reached which defines the maximum available information in the E-mode component of the $\\xi_\\pm$. We show that this limit is reached the earlier (i.e., for a smaller number of modes considered) the narrower the angular range is over which $\\xi_\\pm$ are measured, and it is reached much earlier for logarithmic weight functions. For example, for $\\xi_\\pm$ on the interval $1'\\le \\vt\\le 400'$, the asymptotic limit for the parameter pair $(\\Omega_{\\rm m},\\sigma_8)$ is reached for $\\sim 25$ modes in the linear case, but already for 5 modes in the logarithmic case. The COSEBIs form a natural discrete set of quantities, which we suggest as method of choice in future cosmic shear likelihood analyses.}{} ", "introduction": " ", "conclusions": "We have defined pure E- and B-mode cosmic shear measures from correlation functions over a finite interval $\\tmin\\le \\vt\\le \\tmax$. These are complete orthonormal sets of such measures, implying that they contain all cosmic shear information in the two-point correlation functions which can be uniquely split into E- and B-modes. For these COSEBIs, we have calculated their relation to the underlying power spectrum and their covariance matrix. Two different sets of COSEBIs have been explicitly constructed, those with weight functions which are polynomials in the angular scale, and those with polynomial weight functions in the logarithm of the angular scale. For the former case, analytic expressions were obtained for all orders, whereas in the logarithmic case, a linear system of equations needs to be solved numerically. \\subsection{Advantages of the COSEBIs} Comparing the COSEBIs with earlier cosmic shear measures, we point out a number of advantages. First, using the correlation functions themselves does not provide an E-/B-mode separation. The construction of E-/B-mode correlation functions as described in \\citep{cnp02} requires knowledge of the correlation functions over an infinite angular range, and is therefore not applicable in practice (extrapolating to infinite separation using fiducial cosmological models corresponds to `inventing data', and implicitly assumes that there are no long-range B-modes). In fact, the generalization of pure E-/B-mode correlation functions based on data over a finite angular range has been derived here (see Sect.\\ts\\ref{sc:EB-decomp} and Appendix\\ts\\ref{sc:EBmodeCFs}); however, we expect these to be of limited use in practice. Whereas the aperture mass dispersion \\citep{swj98} provides a clean separation into E- and B-modes \\citep{cnp02,svm02}, it requires the knowledge of the correlation function to arbitrarily small angular separation. There are at least two aspects which render this impractical: first, galaxy images need a minimum separation for their shapes to be measurable. Second, on very small scales baryonic effects will affect the power spectrum and render model predictions very uncertain. The inevitable bias of the aperture mass dispersion \\citep{kse06} motivated the ring statistics \\citep{sck07}. The latter removes the bias, depends only on the correlation function over a finite interval, and has potentially higher sensitivity to cosmological parameters \\citep[][FK10]{esk09}. However, the weight function of the ring statistics is largely arbitrary. The COSEBIs contain all available mode-separable information from the correlation functions on a finite interval, and are therefore guaranteed to provide highest sensitivity to cosmological parameters. Furthermore, they form a discrete set of measures, whereas the other cosmic shear statistics include a somewhat arbitrary grid of variables, like the outer scale of the ring statistics: if the grid is too coarse, information gets lost, whereas a finer grid renders the measures largely redundant, implying large and significantly non-diagonal covariances. In contrast, the discreteness of COSEBIs leaves no freedom, and for the linear weight functions, the covariances have a narrow band structure. The information clearly saturates after a number of modes, and this number is surprisingly small for the logarithmic weight function. Therefore, determining covariance matrices from numerical simulations \\citep[as was done for the COSMOS analysis of][]{shj09} appears considerably simpler than for other cosmic shear measurements, which is particularly true for an unbiased estimate of their inverse \\citep[see][for a discussion of this point]{hss07}. Based on these properties of the COSEBIs, we would like to advertise them as the method of choice for future cosmic shear analyses. \\subsection{Generalizations} In case photometric redshift information of the lensed galaxies is available and several source populations can be defined based on their redshift estimates, the COSEBIs can be generalized to a tomographic version. Furthermore, under the same assumption, intrinsic alignment effects between the tidal gravitational field and the intrinsic galaxy orientation \\citep[e.g.,][]{ckb01,cnp01,jin02,his04} can be filtered out by properly choosing redshift-dependent weight functions, such as to avoid physically close pairs of galaxies \\citep{kis02,kis03,heh03} or make use of the specific redshift dependence of the shear-intrinsic alignments \\citep{brk07,jos08,jos09}, possibly in combination with other data \\citep{job09}. We expect that these generalizations of the COSEBIs provide no real difficulties. It would be desirable to obtain a similar measure for third-order cosmic shear statistics, i.e., one that provides clear E-/B-mode separations from three-point correlation functions measured over a finite interval. Up to now, the aperture statistics is the only known such measure \\citep{jbj04,skl05}; however, similar to the case of the aperture dispersion, third-order aperture statistics requires the correlation functions to be measured down to arbitrarily small separations. A generalization of the COSEBIs to third order seems challening -- not only because of the higher number of independent variables (the three-point correlation functions depend on three variables) and the larger number of modes (one pure E-mode, one mixed E/B-mode, and two further modes which are not invariant under parity transformation), but also because of the more complicated relation between correlation functions and the bispectra \\citep{skl05}. Thus, even the analogue of the starting point of the current paper -- Eqs.\\ts(\\ref{eq:EBmodes}) and (\\ref{eq:Tplusminus}) -- is not yet known for the third-order case." }, "1002/1002.0596_arXiv.txt": { "abstract": "The unprecedented spatial and spectral resolutions of \\ch\\ have revolutionized our view of the X-ray emission from supernova remnants. The excellent data sets accumulated on young, ejecta dominated objects like Cas A or Tycho present a unique opportunity to study at the same time the chemical and physical structure of the explosion debris and the characteristics of the circumstellar medium sculpted by the progenitor before the explosion. Supernova remnants can thus put strong constraints on fundamental aspects of both supernova explosion physics and stellar evolution scenarios for supernova progenitors. This view of the supernova phenomenon is completely independent of, and complementary to, the study of distant extragalactic supernovae at optical wavelenghts. The calibration of these two techniques has recently become possible thanks to the detection and spectroscopic follow-up of supernova light echoes. In this paper, I will review the most relevant results on supernova remnants obtained during the first decade of \\ch, and the impact that these results have had on open issues in supernova research. ", "introduction": "\\subsection{The Power of the Data} The excellent quality of the data generated by \\ch\\ for bright SNRs in our Galaxy is showcased by the false color images that have become perhaps the mission's most recognizable and widespread visual result (Figure 1). For a SNR with an angular diameter of $6 ^\\prime$ like Cas A, \\ch\\ can resolve more than $10^{5}$ individual regions. Deep exposures of bright objects usually provide enough photon statistics to obtain a good spectrum from most of these regions at the moderate spectral resolution of the ACIS CCD detectors ($E / \\Delta E \\approx 10-60$). This usually allows to detect K-shell emission from abundant elements like O, Si, S, Ar, Ca, and Fe. \\begin{figure*}[t] \\begin{center} \\includegraphics[width=\\textwidth]{tycho_kepler_casa.eps} \\caption{Three-color images generated from deep \\ch\\ exposures of the Tycho (left), Kepler (center), and Cas A (right) SNRs. The details vary for each image, but red usually corresponds to low energy X-rays around the Fe-L complex ($\\sim$ 1 keV and below), green to mid-energy X-rays around the Si K blend ($\\sim$ 2 keV), and blue to high energy X-rays in the 4-6 keV continuum bewteen the Ca K and Fe K blends. Images are not to scale: Tycho is $\\sim8\\, ^\\prime$ in diameter, Kepler is $\\sim4\\, ^\\prime$, and Cas A is $\\sim6\\, ^\\prime$. Total exposure times are $150$, $750$, and $1000$ ks. Images courtesy of the \\ch\\ X-ray Center; data originally published in \\cite{warren05:Tycho}, \\cite{reynolds07:kepler}, and \\cite{hwang04:CasA_VLP}.}\\label{fig-1} \\end{center} \\end{figure*} \\subsection{Basic Concepts of NEI Plasmas} The density of the plasma inside SNRs is low enough for the ages of young objects like Cas A or Tycho to be smaller than the ionization equilibrium timescale. The X-ray emitting plasma heated by the shocks is therefore in a state of nonequilibrium ionization, or NEI \\cite{itoh77:SNRs_NEI}. This means that the ionization state of any given fluid element is determined not only by its electron temperature $T_{e}$, but also by its ionization timescale $n_{e}t$, where $n_{e}$ is the electron number density and $t$ is the time since shock passage. The thermal X-ray spectrum from a young SNR is thus intimately related to its dynamic evolution through the individual densities and shock passage times of each fluid element. This has important implications for the ejecta emission, because of the large differences in chemical composition across the SN debris, and the fact that the electron pool is completely dominated by the contributions from high-Z elements, making $n_{e}$ a strong function of the ionization state \\cite{hamilton84:ejecta}. Under these circumstances, the quantitative analysis of X-ray spectra from young SNRs becomes a challenging endeavor. In order to derive magnitudes that are relevant to SN physics, like the kinetic energy $E_{k}$ or the ejected mass of each chemical element, it is necessary to build a hydrodynamic model of the entire SNR. One must know when each fluid element was shocked, what its chemical composition is, how much of the SN ejecta is still unshocked at the present time, etc. Individual fitting of each magnitude becomes impractical, because they are all related to each other, and in order to understand the X-ray spectrum of a particular SNR, one has to understand of the object as a whole. Further complications stem from two sources. One is a technical, but important, problem: the uneven quality of atomic data in X-ray emission codes for NEI plasmas \\cite{borkowski01:sedov,badenes06:tycho}. The other is of a more fundamental nature: the large uncertainties in the physics of collisionless shocks, in particular the amount of ion-electron temperature equilibration at the shock transition \\cite{ghavamian07:shock_equilibration,heng09:balmer_shocks} and the impact of cosmic ray acceleration on the dynamics of the plasma \\cite{decourchelle00:cr-thermalxray}. Most of the analysis of the X-ray emission from SNRs is done following one of two radically different approaches. The first approach is to forgo the complex interaction between SNR dynamics and X-ray spectrum, and fit the emission of the entire SNR or individual regions using ready-made spectral models. The most widely used are plane-parallel shock NEI models \\cite{hughes00:E0102,borkowski01:sedov}, which fit the spectra by varying $T_{e}$, some parameter related to $n_{e}t$, and a set of chemical abundances. This has the advantages of simplicity and flexibility, but the number of free parameters is large, the quality of the fits is often poor, and the results are hard to interpret in the framework of SN physics and progenitor scenarios. The other approach is to model the full HD evolution of the SNR \\textit{ab initio}, starting from a grid of SN explosion models and ISM or CSM configurations, calculate the NEI processes in the shocked plasma, and produce a set of synthetic X-ray spectra that can then be compared to the observations. This approach is less flexible, and usually does not allow for spectral fits in the usual sense (i.e., based on the $\\chi^2$ statistic), but it is considerably more powerful in that the observations and the physical scenarios that are being tested can be compared in a more direct way. The first HD+NEI models for SNRs were built to analyze the data from early X-ray missions like \\textit{Einstein} and \\textit{EXOSAT} (e.g, \\cite{hughes85:nei,hamilton86:SN1006,borkowski94:kepler}). Modern efforts produced for \\ch\\ and \\xmm\\ data can be found in \\cite{badenes03:xray,badenes05:xray} and \\cite{sorokina04:typeIasnrs}. ", "conclusions": "This brief review on the X-ray observations of SNRs has been written with two main goals. The first is to illustrate the power of SNR studies as probes of the SN phenomenon. The unique capabilities of modern X-ray satellites like \\ch\\ offer a radically different view of the explosions and their progenitors, a vision that is independent of, and complementary to, the conventional studies of extragalactic SNe at optical wavelengths. The X-ray data sets assembled for well-known objects like Tycho or Cas A will be a lasting legacy of the \\ch\\ mission, and they represent the most detailed picture of the structure of SN ejecta currently available at any wavelength. The recent discovery that the statistical properties of the ejecta emission resolved by \\ch\\ can be used to distinguish CC from Type Ia SNRs \\cite{lopez09:typing_SNRs} is a powerful example of the potential of these data. It is therefore crucial that the campaign of X-ray observations of SNRs with \\ch\\ and other satellites continue, and that deep exposures be completed for all the objects that do not have them. The discovery of new young SNRs, while difficult, is definitely possible, as illustrated by the $\\sim100$ yr old Galactic SNR G1.9$+$0.3 \\cite{reynolds09:G1.9+0.3}, and should be pursued in parallel to the study of well-known ones. The second goal of this review is to emphasize the importance of careful modeling for the analysis of the X-ray spectra of SNRs. Because of the NEI character of the shocked plasma in SNRs, it is extremely difficult to interpret their X-ray spectra in terms of magnitudes that can be used effectively to constrain explosion physics and progenitor scenarios. A robust, quantitative analysis usually requires putting together a dynamical model for the SNR - in other words, the X-ray emission of a SNR cannot be interpreted without understanding the object as a whole, at least to some degree. In this context, one-dimensional HD+NEI models have been successful in recovering the fundamental properties of SN explosions from the spatially integrated X-ray spectra of Type Ia SNRs, and validation of the technique through the spectroscopy of SN light echoes has been possible in a few cases. These results are certainly encouraging, but much remains to be done before the full potential of the X-ray observations of SNRs is realized. In order to take advantage of the spatially resolved spectroscopy capabilities of \\ch, it is necessary to build multi-dimensional HD+NEI SNR models. This line of research has great promise as a benchmark for multi-dimensional SN explosion models, which hold the key to fundamental processes like the physical mechanism responsible for CC SNe \\cite{janka07:CCSN_Review} or the origin of the diversity within Type Ia SNe \\cite{kasen09:SNIa_diversity}. Present and future X-ray observations of SNRs present many opportunities for SN research. These opportunities are not devoid of challenges, but the excellent quality of the data obtained with \\ch\\ and other satellites fully warrant the effort required to meet them. Because of the unique view of the SN phenomenon that they offer, SNRs should play a central role in shaping our understanding of SN explosions." }, "1002/1002.0844_arXiv.txt": { "abstract": "The baryon content of high-density regions in the universe is relevant to two critical unanswered questions: the workings of nurture effects on galaxies and the whereabouts of the missing baryons. In this paper, we analyze the distribution of dark matter and semianalytical galaxies in the Millennium Simulation to investigate these problems. Applying the same density field reconstruction schemes as used for the overall matter distribution to the matter locked in halos we study the mass contribution of halos to the total mass budget at various background field densities, i.e., the conditional halo mass function. In this context, we present a simple fitting formula for the cumulative mass function accurate to $\\lesssim5\\%$ for halo masses between $10^{10}$ and $10^{15}\\hMsol$. We find that in dense environments the halo mass function becomes top heavy and present corresponding fitting formulae for different redshifts. We demonstrate that the major fraction of matter in high-density fields is associated with galaxy groups. Since current X-ray surveys are able to nearly recover the universal baryon fraction within groups, our results indicate that the major part of the so-far undetected warm--hot intergalactic medium resides in low-density regions. Similarly, we show that the differences in galaxy mass functions with environment seen in observed and simulated data stem predominantly from differences in the mass distribution of halos. In particular, the hump in the galaxy mass function is associated with the central group galaxies, and the bimodality observed in the galaxy mass function is therefore interpreted as that of central galaxies versus satellites. ", "introduction": "The environmental dependence of galaxy properties, such as broadband color, star formation, and stellar mass, is a well-known effect in the local universe \\citep{Dressler-80}. In this context environment means an estimate of the smoothed density filed at a given location. Of particular interest for the present study is the dependence of the galaxy stellar mass function (GMF) on environment. Recent comprehensive galaxy redshift surveys have led to intensive studies in this field. For instance, \\cite{Mo-04} model the dependence of the luminosity function on the large-scale environment based on mock catalogs for the the two-degree Field Galaxy Redshift Survey \\citep[2dFGRS;][]{Colless-01} and \\cite{Baldry-06} investigate the GMF as a function of environment in the nearby universe based on the Sloan Digital Sky Survey \\citep[SDSS;][]{York-00}. To uncover evolutionary effects, surveys spanning a larger redshift range have been used: the study by \\cite{Bundy-06} is based on the DEEP2 Galaxy Redshift Survey ($0.4\\leq z \\leq 1.4$); \\cite{Pannella-09} use the COSMOS survey; \\cite{Bolzonella-09} employ the zCOSMOS survey in the redshift range $0 -14.5$ dwarf irregular galaxies. We start with the measured (beam corrected) distribution of apparent axial ratios in the HI 21cm images of dwarf irregular galaxies observed as part of the Faint Irregular Galaxy GMRT Survey (FIGGS). Assuming that the discs can be approximated as oblate spheroids, the intrinsic axial ratio distribution can be obtained from the observed apparent axial ratio distribution. We use a couple of methods to do this, and our final results are based on using Lucy's deconvolution algorithm. This method is constrained to produce physically plausible distributions, and also has the added advantage of allowing for observational errors to be accounted for. While one might a priori expect that gas discs would be thin (because collisions between gas clouds would cause them to quickly settle down to a thin disc), we find that the HI discs of faint dwarf irregulars are quite thick, with mean axial ratio $ \\sim 0.6$. While this is substantially larger than the typical value of $\\sim 0.2$ for the {\\it stellar} discs of large spiral galaxies, it is consistent with the much larger ratio of velocity dispersion to rotational velocity ($\\sigma/v_c$) in dwarf galaxy HI discs as compared to that in spiral galaxies. Our findings have implications for studies of the mass distribution and the Tully - Fisher relation for faint dwarf irregular galaxies, where it is often assumed that the gas is in a thin disc. ", "introduction": "\\label{sec:int} The intrinsic shape of galaxies is interesting from a variety of perspectives. For a given galaxy the shape should be consistent with the dynamical model of the galaxy, while, for a sample of galaxies, one would expect that a correct evolutionary model should be able to reproduce the observed distribution of shapes. It is generally assumed that disc galaxies can be approximated to be oblate spheroids (e.g. \\cite{hubble26,sandage70,ryden06}). If one further assumes that the galaxies have a well defined mean axial ratio (q$_0$), then the observed axial ratio can be used to determine the inclination of the disc. In turn, the inclination is a crucial input in dynamical modeling (e.g. for the mass and structure of the dark matter halo), studying the Tully-Fisher relation etc. The observed shape of a galaxy differs from the intrinsic shape because of projection effects. If one has a sample of galaxies drawn from a population with a well defined intrinsic axial ratio distribution and with random orientations with respect to the earth, then one can determine the distribution of intrinsic axial ratios from the observed axial ratio distribution \\citep[for e.g.~][]{noe79,bin81,lam92,ryden06}. It is worth noting that most of these studies have focused on large galaxies, and that there have been relatively few that focused on dwarfs. For bright spiral galaxies, q$_0$ is often taken to be $\\sim 0.2$ \\citep[see e.g.~][]{haynes84,verheijen01}. It has also long been appreciated that the axial ratio is a function of Hubble type. For example, \\cite{heidmann72} found that while discs get thinner as one goes from galaxies of morphological type Sa to Sd, there is a rapid increase in disc thickness as one goes from Sd to dwarf Irregular galaxies. Similarly, \\cite{sta92} found that dwarf galaxies from the UGC catalog have q$_0 \\sim 0.5$. Axial ratio is also a function of the wavelength of observation. For example, \\citet{ryden06} showed that older populations as traced by redder stars have thicker ratios than the corresponding B band disc. But all of these studies refer to the {\\it stellar} discs of the galaxies. Due to collisions between gas clouds, one would expect that the gas discs would be intrinsically quite thin, for example, for our Galaxy, the scale height of the middle disc is only $\\sim 300$~pc. In extremely gas rich dwarf galaxies one might then expect that the gas discs are relatively thin, even though the stellar disc may have a large axial ratio. We present here a study of the axial ratio distribution of the {\\it HI discs} of extremely faint dwarf irregular galaxies. We select a sample of galaxies for which HI synthesis observations are available from the Faint Irregular Galaxy GMRT Survey (FIGGS, \\cite{ay08}). The sample selection and analysis methods used are presented in section~\\ref{sec:sample}, while the results are discussed in section~\\ref{sec:dis}. To the best of our knowledge, this is the first such study of the intrinsic shape of the HI discs of galaxies. ", "conclusions": "\\label{sec:dis} Eqn.~\\ref{eqn:i1} (and hence results obtained from inverting it also) applies only for a sample of randomly oriented oblate spheroids. The basic assumption hence is that the galaxy sample we are working with has an unbiased distribution of inclination angles. The selection criteria for the FIGGS sample (from which the current sample is drawn) includes a requirement that the optical major axis of the galaxy be $> 1^{'}$. Dwarf galaxies are generally dust poor \\citep[for eg. see ][]{wal07,gal09}, and hence to a good approximation have optically thin discs. Highly inclined galaxies will hence be over represented in a diameter limited sample, i.e. our sample is biased towards edge on discs. This means that the true mean intrinsic axial ratio would be even larger than what we have estimated above. It is worth noting that as the intrinsic axial ratio gets closer to 1.0, the magnitude of this bias decreases, and hence the bias in our estimate should not be large. The mean intrinsic axial ratio that we obtain, viz. $ \\sim 0.57$ is substantially larger than the value of $0.2$ usually adopted for the {\\it stellar } discs of large spiral galaxies. The value from our sample is more than twice as large as older measurements of $$ in stellar discs of Magellanic irregular galaxies by \\citet{heidmann72} (ranging from 0.20 to 0.24). Consistent with this, our sample contains no very flat galaxies. In fact, as can be seen from Fig.~\\ref{fig:bind} there are no galaxies with apparent axial ratio $<0.2$ in the sample. Further, if we look at the distribution of the observed axial ratio of different classes of galaxies in the Automated Photographic Measuring (APM) survey as in \\citet{lam92}, then our histogram resembles those for ellipticals and S0s more closely than that for spirals. This is another qualitative indication that the underlying intrinsic distribution of axial ratios has a higher mean than is typical for spirals. Interestingly, the value of $$ we obtain matches well with what \\citet{sta92,binggeli95} derived for the {\\it stellar} discs in dwarfs. The thickness of the {\\it gas } discs of dwarf galaxies is contrary to what one might have naively expected for a gas disc, since, in general, collisions between gas clouds should cause them to quickly settle into a thin disc. However, this large axial ratio is probably consistent with the large gas dispersion in comparison to the rotational velocity observed in dwarf galaxies. For example, \\citet{kau07} did particle hydrodynamic simulations to show that dwarf galaxies with rotational velocities $\\sim$40 \\kms~did not originate as thin discs but thick systems. This still leaves open the question of where the large velocity dispersion comes from. \\cite{dut09} find good evidence for a scale free power spectrum of HI fluctuations in dwarf galaxies, consisent with what would be expected from a turbulent medium. Interestingly, they also find that that the dwarfs must have relatively thick gas discs, similar to the conclusions reached here. Assuming that the origin of the velocity dispersion is turbulent motions in the ISM, the timescale for dissipation is given by $\\tau \\sim L/v_{turb} \\sim L/\\sigma$ \\citep[e.g.]{shu87}. The total turbulent energy is $E_{turb} \\sim 1/2 M_{HI} \\sigma^2$. For our sample galaxies, the typical HI mass is $\\sim 3 \\times 10^7$~M$_\\odot$, while the length scale $L \\sim 1$kpc. The rate of turbulent energy dissipation is hence $\\sim 10^{44}$erg/yr. On the other hand, the star formation rate is $\\sim 10^{-3}$ M$_\\odot$/yr \\citep{roychowdhury09}, for which the expected supernova rate for a Salpeter IMF is $\\sim 7 \\times 10^{-6}$/yr \\citep{bin01}. Assuming that each supernova explosion deposits $\\sim 10^{51}$ ergs of energy into the ISM, the energy input from star formation is $\\sim 10^{46}$ ergs/yr, more than sufficient to balance the turbulent energy loss. Thus, star formation driven turbulence in the ISM is a plausible cause for the thick gas discs that we observe. We have assumed through out that the HI discs of dwarf galaxies can be approximated as oblate spheroids. On the other hand, in Sec.~\\ref{sec:sample} we saw that the inversion based on the best fit polynomial to the observed histogram of axial ratios actually gives unphysical results, i.e. that the derived intrinsic axial ratio distribution \\psiq\\ becomes negative. \\cite{lam92} found a similar pattern for the axial ratio distribution of spiral and S0 galaxies in the APM catalog, and hence relaxed the assumption that the discs are oblate spheroids. Adequate fits to their data could be obtained assuming that the galaxies have a triaxial shape. Similarly, \\cite{ryden06} from a study of large galaxies in the 2MASS catalog concluded that spiral galaxies are mildly triaxial. The fact that the polynomial approximation gives unphysical results for our sample also suggests that triaxial models may provide a better fit. On the other hand, the Lucy deconvolution gives a physically plausible (as indeed it is designed to) intrinsic axial ratio distribution, that fits the observed data within the error bars. It is worth noting that distribution found by Lucy deconvolution is in excellent agreement with that found by direct inversion of the polynomial fit for the entire range for which the latter is $>0$. Interestingly, if we assume the gas discs to be prolate instead of oblate spheroids, Lucy deconvolution produces an equally acceptable \\phip~ as can be seen from Fig.~\\ref{fig:prophi}. The \\psiq~ obtained from Lucy deconvolution with a prolate spheroid assumption (see Fig.~\\ref{fig:propsi}) indicates that such gas discs should be more cylindrical than spherical. However, the observed kinematics of gas in these galaxies \\citep[see ][etc.]{bch03,bc03,bc04a,bc04b,bck05,bcks05,beg06,ay08} shows that the gas has a significant rotational support. As discussed earlier, a gravitationally bound rotating disc of gas forms an oblate and not a prolate spheroid. Thus, although from the axial ratio data alone, one cannot distinguish between a prolate and oblate shapes, in conjuction with the kinematical data, it is clear that the gas is distributed in the form of a thick disc. \\begin{figure} \\psfig{file=prophi.ps,height=3.4truein, angle=270} \\caption{Reconstructed axial ratio distribution ($\\Phi(p)$) assuming the gas discs to be prolate spheroids, is shown as boxes, which follows the best fit polynomial to the binned data (solid line) closely. Binned data with errorbars and the reconstructed distribution $\\Phi(p)$ obtained using the oblate spheroid assumption (dashed line) is shown for comparison. All values have been normalized to unity.} \\label{fig:prophi} \\end{figure} \\begin{figure} \\psfig{file=propsi.ps,height=3.4truein, angle=270} \\caption{The frequency distribution of {\\it intrinsic} axial ratios obtained after Lucy deconvolution. The dashed line indicates \\psiq~ obtained assuming the gas discs to be prolate spheroids, while the solid line shows the earlier \\psiq~ obtained with the oblate spheroid assumption. } \\label{fig:propsi} \\end{figure} In summary, we find that the {\\it gas} discs of dwarf galaxies are relatively thick, in sharp contrast to the gas discs of spiral galaxies. This has implications both for the internal dynamics of the gas, as well as for studies of the mass distribution Tully - Fisher relation in faint dwarfs, in which it is generally assumed that the gas is in a thin disc." }, "1002/1002.1067_arXiv.txt": { "abstract": "{We investigate the properties of bright ($M_{V} \\le -18$) barred and unbarred disks in the Abell 901/902 cluster system at $z\\sim$~0.165 with the STAGES HST ACS survey. To identify and characterize bars, we use ellipse-fitting. We use visual classification, a S{\\'e}rsic cut, and a color cut to select disk galaxies, and find that the latter two methods miss 31\\% and 51\\%, respectively of disk galaxies identified through visual classification. This underscores the importance of carefully selecting the disk sample in cluster environments. However, we find that the global optical bar fraction in the clusters is $\\sim$~30\\% regardless of the method of disk selection. We study the relationship of the optical bar fraction to host galaxy properties, and find that the optical bar fraction depends strongly on the luminosity of the galaxy and whether it hosts a prominent bulge or is bulgeless. Within a given absolute magnitude bin, the optical bar fraction increases for galaxies with no significant bulge component. Within each morphological type bin, the optical bar fraction increases for brighter galaxies. We find no strong trend (variations larger than a factor of 1.3) for the optical bar fraction with local density within the cluster between the core and virial radius ($R\\sim$~0.25 to 1.2 Mpc). We discuss the implications of our results for the evolution of bars and disks in dense environments. ", "introduction": "\\label{intro} Mounting evidence suggests that a dominant fraction of bulges in fairly massive ($M_{\\ast}\\sim 10^{10} - 10^{11}$~M$_{\\odot}$) $z\\sim$~0 galaxies (e.g., Weinzirl et al. 2009) as well as the bulk of the cosmic star formation rate density since $z < $~1 in intermediate and high mass galaxies, are not triggered by ongoing major mergers \\citep{Hammer05, Wolf05, Bell05, JogeeIAU, Jogee08} but are likely related to a combination of minor mergers and internally-driven secular processes Barred disks have been studied extensively in the nearby universe. Quantitative methods show that approximately 45\\% of disk galaxies are barred in the optical at $z\\sim$~0 \\citep{MJ07, Reese07, BJM08, Aguerri08}. Bars drive gas to the centers of galaxies, where powerful starbursts can be ignited \\citep{Schwarz81, KK04, Jogee05, Sheth05}. These starbursts caused by bar-driven inflow may help build up central, high $v/\\sigma$, stellar concentrations called disky bulges or pseudobulges \\citep{Kormendy82, Kormendy93, Jogee05, Weinzirl08}. Bars can also create boxy/peanut bulges through vertical buckling \\citep{Bureau99, Ath05, MVShlos06} and resonance scattering \\citep{Combes81, Combes90}. But what is the relationship between bar-driven secular evolution and environmental effects? Little is known about barred galaxies in dense environments, as the situation is complicated by the relative importance of processes such as ram pressure stripping, galaxy tidal interactions, mergers, and galaxy harassment. In addition, further complications are introduced by the fact that the bar fraction and properties in clusters depend on the epoch of bar formation and the evolutionary history of clusters. We explore these questions using the Space Telescope Abell 901/902 Galaxy Evolution Survey (STAGES; \\citealp{Gray08}). ", "conclusions": "" }, "1002/1002.2472_arXiv.txt": { "abstract": "The energy spectra of ultra-high energy cosmic rays (CRs) measured with giant extensive air shower (EAS) arrays exhibit discrepancies between the flux intensities and/or estimated CR energies exceeding experimental errors. The well-known intensity correction factor due to the dispersion of the measured quantity in the presence of a rapidly falling energy spectrum is insufficient to explain the divergence. Another source of systematic energy determination error is proposed concerning the charged particle density measured with the surface arrays, which arises due to simplifications (namely, the superposition approximation) in nucleus-nucleus interaction description applied to the shower modeling. Making use of the essential correction factors results in congruous CR energy spectra within experimental errors. Residual differences in the energy scales of giant arrays can be attributed to the actual overall accuracy of the EAS detection technique used. CR acceleration and propagation model simulations using the dip and ankle scenarios of the transition from galactic to extragalactic CR components are in agreement with the combined energy spectrum observed with EAS arrays. ", "introduction": "Ultra-high energy cosmic rays (UHECRs) are presently measured using a number of giant extensive air shower (EAS) arrays. The observed energy spectrum exhibits the cutoff predicted by \\citet{G} and \\citet{ZK} (GZK) and `Ankle' features at energies of approximately $4\\times10^{19}$ eV and $5\\times10^{18}$ eV \\citep{Fksm}. However, there are essential discrepancies between the flux intensities and/or the estimated energies of the initial cosmic rays (CRs) generating EASs detected by different arrays. These discrepancies exceed instrumental errors and it makes it difficult to decide for the validity of the results obtained. The current paper presents an analysis of the sources of these discrepancies. It is shown by i) correcting the CR intensity due to instrumental errors and power law spectrum and ii) taking into account model uncertainty in the estimation of EAS initial nucleus energy that the observed UHECR energy spectra appear to be congruent. Residual differences in UHECR energy scales of arrays can be attributed to the actual overall accuracy of the EAS detection technique. ", "conclusions": "It has been noted previously that UHECR energy spectra measured with giant EAS arrays agree with each other if the energy scales are adjusted (\\cite{BW,DAN,Brznsk}, and other papers cited therein). However, the energy correction factors needed to merge the spectra exceed experimental errors. For example, \\citet{Brznsk} used values 1.2, 0.75, 0.83, and 0.625 to shift PAO, AGASA, Akeno, and Yakutsk data to those of HiRes. Moreover, instrumental and modeling errors are believed to equiprobably increase or decrease the estimated energy. Instead, as \\citet{Wtsn} concluded when comparing integral calibrated fluxes of UHECRs from giant arrays, $S_{600}$ measurements and fluorescence measurements of EAS assemble in two separate groups of data (Volcano Ranch, Haverah Park, AGASA vs PAO, HiRes results). \\begin{figure}[t] \\includegraphics[width=\\columnwidth]{SpectraREA.eps} \\caption{The observed spectra with energy scales adjusted. Correction factors are taken from Table~\\ref{Table:RA}. CR intensities and data symbols are the same as in Fig.~\\ref{fig:SpectraRJ} and ~\\ref{fig:SpectraRA}. Model calculation results: full line \\citep{Dip}; dot and dash lines \\citep{AGN}; dash-dot line \\citep{Wbg}.} \\label{fig:SpectraRE} \\end{figure} Considerations in this paper insist that, besides the `Moscow correction factor'\\footnote{derived by \\cite{Ztspn}, \\cite{Mrzn}, and \\cite{NNK}}, which reduces all CR intensities measured with EAS arrays, there is another factor that reduces the estimated energy of primary nuclei diminishing a difference between two groups of data (Fig.~\\ref{fig:SpectraRA}). The remaining spread of $\\hat{E}$ can be attributed to the real accuracy of the EAS measurement technique: instrumental errors + model uncertainty. Two variants of the energy correction factors for EAS arrays are given in Table \\ref{Table:RA} for the two extreme cases of the primary nuclei (H,Fe). Resultant spectra with adjusted energy scales are shown in Fig.~\\ref{fig:SpectraRE}. Energy correction factors are assumed constant at energies above $1$ EeV, and the spectra are shifted to a sample mean. Having an indefinite mass of EAS primary nuclei and model uncertainties to assign the energy correction factors, we can find an interval of values where $R_E$ would be. A sample of six experiments (Figs.~\\ref{fig:SpectraRJ}, \\ref{fig:SpectraRA}) estimates that the average energy determination error, $\\delta \\hat{E}/\\hat{E}$, inherent to giant EAS arrays, is within an (8,19)\\% interval. Now we can compare the measured energy spectra with model simulations of extragalactic (EG) CRs predominating over the Galactic (G) component. The dip (described as `bump' in early measurements) observed in the CR spectrum~\\citep{Bump1,Bump2} was explained by \\citet{Dip0}\\footnote{it has been updated a number of times since; the latest is made by \\citet{Dip}}, who studied the Bethe--Heitler pair production of protons from distant sources on the cosmic microwave background. The authors conclude that the dip shape is formed by the universal modification factor and is independent of UHECR propagation details. The only phenomenon known to modify the dip is the presence of a significant fraction of nuclei in the primary beam. In Fig.~\\ref{fig:SpectraRE} the spectrum calculated in the dip scenario \\citep{Dip} with protons accelerated in EG sources\\footnote{acceleration spectrum index $\\gamma=2.7$} is shown. Only the source luminosity is fitted to the experimental data. The position on the energy scale and shape of the dip agree with observed spectra within experimental errors, as well as the GZK effect. The excess-over-GZK flux observed by the AGASA array is considered to be the result of the primary-energy overestimation in inclined showers of the highest energies \\citep{Cpdvll}. \\begin{table}[t]\\begin{center} \\caption{Energy correction factors for EAS arrays. \\label{Table:RA}} \\begin{tabular}{lccccc}&&&&&\\\\ \\hline Array & AGASA & HP & Yakutsk & PAO & HiRes/TA\\\\ \\hline $R_E^{Fe}$ & 0.96 & 0.98 & 0.88 & 1.15 & 1.01\\\\ $R_E^{H}$ & 0.81 & 1.01 & 0.71 & 1.29 & 1.09\\\\ \\hline \\end{tabular}\\end{center} \\end{table} \\citet{Wbg} have argued that the ankle in the spectrum marks the transition from G to EG components of CRs. This phenomenological `ankle scenario' uses a sum of power law G and EG spectra with slopes differing by $\\Delta\\gamma\\sim1.8$. The result is shown in Fig.~\\ref{fig:SpectraRE} by the dash-dot line. The authors emphasize the sharpness of the ankle observed, $d^2\\ln JE^3/d\\ln^2E$, which is consistent with the ankle scenario, while in the dip scenario the sharpness is of insufficient magnitude, especially in the case of a large nuclei-to-proton ratio in the primary beam. The energy spectrum of UHECRs produced at the shock created by the expanding cocoons around active galactic nuclei combined with the G component of CRs produced in supernova remnants was calculated by \\citet{AGN}. Expected CR composition shows an increase of $\\overline{A}$ at $E_0\\sim0.1$ EeV owing to the G component, and a second one at energy $\\sim10$ EeV, produced by nonrelativistic cocoon shocks. Calculated UHECR intensity \\citep{AGN} as a function of energy is shown in Fig.~\\ref{fig:SpectraRE} for the two cases: dip (dots) and ankle (dashed line) scenarios of the G to EG transition. Experimental errors are too large to distinguish between alternative scenarios." }, "1002/1002.0707_arXiv.txt": { "abstract": "{Massive, eclipsing, double-lined, spectroscopic binaries are not common but are necessary to understand the evolution of massive stars as they are the only direct way to determine stellar masses. They are also the progenitors of energetic phenomena such as X-ray binaries and $\\gamma$-ray bursts.} {We present a photometric and spectroscopic analysis of the candidate binary system Cyg OB2-B17 to show that it is indeed a massive evolved binary.} {We utilise $V$ band and white-light photometry to obtain a light curve and period of the system, and spectra at different resolutions to calculate preliminary orbital parameters and spectral classes for the components.} {Our results suggest that B17 is an eclipsing, double-lined, spectroscopic binary with a period of $4.0217\\pm0.0004$ days, with two massive evolved components with preliminary classifications of O7 and O9 supergiants. The radial velocity and light curves are consistent with a massive binary containing components with similar luminosities, and in turn with the preliminary spectral types and age of the association.} {} ", "introduction": "O stars are amongst the most massive and intrinsically luminous stellar objects found in galaxies. Since they are the only direct way to measure the masses and radii of stars, binaries are the perfect testbeds for studying the physical properties and evolution of such stars. Unfortunately, as reported by \\cite{bon08}, less than 20 O stars have accurate ($\\leq$10\\%) dynamical mass estimates. Because of the scarcity of known massive, eclipsing, double-lined, spectroscopic binaries - and hence dynamical mass estimates for stars at different evolutionary states \\citep[e.g.,][]{gie02} - the mass luminosity relation and theoretical evolutionary tracks of massive stars ($M\\geq20M_\\odot$) are currently poorly constrained by observations. In order to address this shortfall, much effort has been expended to identify further examples, utilising photometric and spectroscopic observations of young massive clusters such as the Arches \\citep{mar08}, Quintuplet \\citep{fig99}, and Westerlund 1 (\\citealt{cla05}; \\citealt{rit09}). These observations indicate that the binary fraction is potentially very high\\footnote{A large amount of observations and patience are necessary to determine the true binarity of a sample of stars in an open cluster as observed in \\cite{san08}}; \\cite{kob07} inferred it to be $\\geq$ 80\\% in Cygnus OB2 (Cyg OB2), while \\cite{cla08} estimate the binary fraction of WR stars in Westerlund 1 to be $\\geq$70\\%. Currently the stars with the highest dynamical mass estimates are the twin WN6ha components of WR20a within Westerlund 2 (83$M_\\odot$+82$M_\\odot$; \\citealt{rau04}, \\citealt{bon04}) and the newly discovered WN6ha binary NGC3603-A1 (116$\\pm$31$M_\\odot$ + 89$\\pm$16$M_\\odot$; \\citealt{sch08}). This in turn has implications for the determination of the Initial Mass Function (IMF) for the clusters in question and, by extension, the empirically determined maximum mass possible for a star \\citep[cf.][]{fig05}. Moreover, massive close binaries are the progenitors of such diverse energetic phenomena as supernovae, $\\gamma$-ray bursts and X-ray binaries \\citep{rib06}. Clearly the properties of the progenitor binary population must be known to constrain their formation channels. Cyg OB2 is one of the most massive and richest associations in the Galaxy. It is $\\sim$2 Myr old and 1.8 kpc away \\citep{kim07}. Containing at least 60-70 O-type stars \\citep{neg08}, its proximity and accessibility to optical studies have made it the focus of numerous observational campaigns to determine the properties of its massive stellar population \\citep[e.g.][]{mas91, kno00, han03, com02, kim07, kim08}. Cyg OB2 B17 (\\citealt{com02}; henceforth B17 and also known as V1827Cyg, 2MASS J20302730+4113253, NSVS 5738756; $\\alpha_{2000}$=20$^{h}$30$^{m}$27.3$^{s}$, $\\delta_{2000}$=+41$^{o}13^{\\prime}$25$^{\\prime\\prime}$; $V$=12.6) is a luminous, variable member of the Cyg OB2 association. \\citet{com02} observed the system as part of a near-infrared spectroscopic survey and found the spectrum presented $Br$$\\gamma$ emission, confirming it to be an evolved massive star. Follow up observations were made by \\cite{neg08}, who found it to stand out from the rest of the members due to its variability and strong emission lines. They classified it as an Ofpe star and suggested it was a strong binary candidate. This paper reports the first results of an extensive multi-epoch photometric and spectroscopic observational campaign on the binary candidate B17. Section~2 gives a description of the photometric and spectroscopic observations. Photometrically, the system was found to be variable and we report the analysis of the light curve in Section~3. More spectroscopic data were obtained permitting preliminary spectral and luminosity classification of the system; this analysis is shown in Section~4, along with descriptions of the long and short timescale variations of the spectra. The light and radial velocity curve modelling is presented in Section~5. A discussion of the system, including its evolutionary status, is found in Section~6 and a summary is presented in Section~7. Note that the central goals of this manuscript are to verify the binary hypothesis and present a preliminary spectral classification. A full analysis of the system, consisting of the deconvolution of an expanded spectral data set and subsequent model atmosphere analysis to determine the fundamental stellar parameters of the system will be presented in a future paper (Stroud et al. in prep). ", "conclusions": "Using photometric and spectroscopic data, we have demonstrated that B17 is an eclipsing, double lined spectroscopic binary comprising two supergiants with preliminary classifications of O7Iaf and O9Iaf. The spectra are highly variable, and with a subset revealing features from both stars, raise the possibility of achieving a dynamical mass determination for both components. Utilising both the photometric lighturve and our limited RV dataset we attempted to determine an initial orbital solution for the binary. Despite the morphology of the lightcurve indicating a semi-contact configuration we were unable to to achieve convergence for such a hypothesis and hence were forced to adopted an over-contact configuration. In the absence of a full RV curve for both system components we were forced to fix the binary mass ratio, and had to include the presence of a star spot to address the observed asymmetries in the lightcurve (which are likely due to the effects of binary mass transfer). However, we were still unable to fully fit both secondary eclipse and the lightcurve between orbital phase $\\sim$0.6-0.9 with such a model; as such we regard the modelling results presented in this work as provisional. We anticipate that a full RV curve for both components will be necessary to obtain more precise parameters for the system; additional data to accomplish this goal are currently being obtained and refined analysis will be presented in a future paper. Nevertheless, the provisional distance calculation of 1.5-1.8 kpc obtained from the lightcurve analysis agrees with previously published values for the distance to Cyg OB2, being inconsistent with the distance of 900-950 pc determined by \\cite{lin09} with the Cyg~OB2 \\#5 light curve analysis. When placed in the HR diagram, B17 appears to be consistent with the age and stellar population of the Cyg OB2 association. Assuming the system avoids merger, it is likely to evolve through an extreme B supergiant/LBV phase into a long period WR+WR binary configuration as mass loss via stellar winds increases the orbital separation. In combination with the recent work of \\cite{kim09} and \\cite{kob07} the results of our analysis provides additional evidence that Cyg~OB2 has a very high fraction of massive binary stars. Such an observational constraint needs to be considered when determining the initial mass function of the association as it may both influence the slope of the relationship and also lead to a population of artificially massive stars, resulting in the inflation of a putative high mass cut-off to the IMF." }, "1002/1002.2534_arXiv.txt": { "abstract": "We report the first two-dimensional stellar population synthesis in the near-infrared of the nuclear region of an active galaxy, namely Mrk\\,1066. We have used integral field spectroscopy with adaptative optics at the Gemini North Telescope to map the to map the age distribution of the stellar population in the inner 300\\,pc at a spatial resolution of 35\\,pc. An old stellar population component (age $\\gtrsim$5\\,Gyr) is dominant within the inner $\\approx$\\,160\\,pc, % which we attribute to the galaxy bulge. Beyond this region, up to the borders of the observation field ($\\sim$\\,300\\,pc), intermediate age components (0.3--0.7\\,Gyr) dominate. We find a spatial correlation between this intermediate age component and a partial ring of low stellar velocity dispersions ($\\sigma_*$). Low-$\\sigma_*$ nuclear rings have been observed in other active galaxies and our result for Mrk\\,1066 suggests that they are formed by intermediate age stars. This age is consistent with an origin for the low-$\\sigma_*$ rings in a past event which triggered an inflow of gas and formed stars which still keep the colder kinematics (as compared to that of the bulge) of the gas from which they have formed. At the nucleus proper we detect, in addition, two unresolved components: a compact infrared source, consistent with an origin in hot dust with mass $\\approx1.9\\times10^{-2}$\\,M$_\\odot$, and a blue featureless power-law continuum, which contributes with only $\\approx$15\\% of the flux at 2.12\\,$\\mu$m. ", "introduction": "Optical spectroscopy on scales of hundreds of parsecs around the nucleus of Seyfert galaxies have shown that in $\\approx$\\,40\\,\\% of them the active galactic nucleus (AGN) and young stars co-exist \\citep[e.g.][]{sb00,sb01,gd01,cid04,asari07,dors08}, % providing support to the so-called AGN-Starburst connection \\citep[e.g.][]{norman88,terlevich90,heckman97,heckman04,n7582}. In particular, these authors have pointed out that the main difference between the stellar population (hereafter SP) of active and non-active galaxies is an excess of mainly intermediate age stars in the former. A similar result has been found in recent SP studies in the near-infrared (hereafter near-IR), using the technique of spectral synthesis \\citep{rogerio09,rogerio07}. Using integrated spectra of the central few hundreds of parsecs in a sample of 24 Seyfert galaxies, these authors have shown that the continuum is dominated by the contribution of intermediate-age stellar population components (SPCs). In addition, they found that the near-IR nuclear spectra of about 50\\,\\% of the Seyfert~1 and $\\sim$20\\,\\% of the Seyfert~2 galaxies show emission from hot dust \\citep{rogerio09,n4151,n7582,ardila06,ardila05}. Using a somewhat distinct technique, \\citet{davies07} obtained near-IR integral field spectroscopy to investigate the circumnuclear star formation in 9 nearby Seyfert galaxies at spatial resolutions of tens of parsecs. They have modelled the \\br\\ equivalent width, supernova rate and mass-to-light ratio to quantify the star formation history in the center of these galaxies using their code {\\sc stars}. They found that the ages of the stars which contribute most to the near-IR continuum lie in the range 10--30\\,Myr, but point out that these ages should be considered only as ``characteristic\", as they have not performed a proper spectral synthesis, arguing that there may be simultaneously two or more SPs that are not coeval \\citep{davies07,davies06}. In this paper we present, for the first time, two-dimensional (hereafter 2D) SP synthesis in the near-IR for the inner hundreds of pc of an active galaxy -- namely Mrk\\,1066 -- using integral field spectroscopy with adaptive optics, which allowed us to derive the contribution of distinct SPCs to the near-IR spectra and map their spatial distributions. This paper is organized as follows. In Sec.~\\ref{data} we describe the observations, data reduction procedures and the spectral synthesis method; in Sec.~\\ref{results} we present our results, which are discussed in Sec.~\\ref{discussion}. The conclusions are presented in Sec.~\\ref{disc}. ", "conclusions": "\\label{disc} The present work reports for the first time spectral synthesis in the near-IR with 2D coverage for the nuclear region of a Seyfert galaxy (Mrk\\,1066) within the inner $\\approx$\\,300\\,pc at a spatial resolution of $\\approx$\\,35\\,pc. We have mapped the distribution of stellar population components of different ages and of their average reddening. The main conclusions of this work are: \\begin{itemize} \\item The age of the dominant stellar population presents spatial variations: the flux and mass contributions within the inner $\\approx$\\,160~pc are dominated by old stars ($t \\ge$5\\,Gyr), while intermediate age stars ($0.3 \\leq t \\leq 0.7$~Gyr) dominate in the circumnuclear region; \\item There is a spatial correlation between the distribution of the intermediate age component and low stellar velocity dispersion values which delineate a partial ring around the nucleus. Similar structures have been found around other active nuclei, and our result for Mrk\\,1066 suggests that these nuclear rings (and in some cases disks) are formed by intermediate age stars. \\item There is an unresolved dusty structure at the nucleus with mass $M_{\\rm HD}\\approx1.9\\times10^{-2}$~M$_\\odot$, which may be the hottest part of the dusty torus postulated by the unified model of AGN and a small contribution from a power-law continuum ($\\approx$15\\,\\% of the flux at 2.12\\,$\\mu$m); \\item The near-IR synthesis seems not to be sensitive to very recent star formation (with $t\\lesssim5\\,$Myr), reinforcing the importance of multi-wavelength stellar population studies of the central region of active galaxies. \\end{itemize}" }, "1002/1002.1701_arXiv.txt": { "abstract": "A non-resonant instability for the amplification of the interstellar magnetic field in young Supernova Remnant (SNR) shocks was predicted by \\citet{Bell:2004p24}, and is thought to be relevant for the acceleration of cosmic ray (CR) particles. For this instability, the CRs streaming ahead of SNR shock fronts drive electromagnetic waves with wavelengths much shorter than the typical CR Larmor radius, by inducing a current parallel to the background magnetic field. We explore the nonlinear regime of the non-resonant mode using Particle-in-Cell (PIC) hybrid simulations, with kinetic ions and fluid electrons, and analyze the saturation mechanism for realistic CR and background plasma parameters. In the linear regime, the observed growth rates and wavelengths match the theoretical predictions; the nonlinear stage of the instability shows a strong reaction of both the background plasma and the CR particles, with the saturation level of the magnetic field varying with the CR parameters. The simulations with CR-to-background density ratios of $n_\\mathrm{CR}/n_\\mathrm{b}=10^{-5}$ reveal the highest magnetic field saturation levels, with energy also being transferred to the background plasma and to the perpendicular velocity components of the CR particles. The results show that amplification factors $>10$ for the magnetic field can be achieved, and suggest that this instability is important for the generation of magnetic field turbulence, and for the acceleration of CR particles. ", "introduction": "Very energetic CRs ($\\sim10^{14}\\,\\mathrm{eV}$ to $10^{15}\\,\\mathrm{eV}$) are thought to be accelerated in SNR shocks. There is direct evidence that electrons are accelerated up to energies of $10^{14}\\,\\mathrm{eV}$ at SNR sites \\citep{KOYAMA:1995p201,Allen:1997p286,tanimori:1998p917,Aharonian:2001p352,Naito:1999p489,Aharonian:1999p497,Berezhko:2003p504,Vink:2003p515}, and the measured power law spectra of the CRs indicates Diffusive Shock Acceleration (DSA) as the most likely mechanism responsible for the acceleration \\citep{Axford:1977p671,Bell:1978p884,Blandford:1978p891}. The acceleration of these particles up to energies of $\\sim10^{15}\\,\\mathrm{eV}$ through the DSA mechanism requires the existence of magnetic fields much stronger than the typical $B_0\\sim3\\,\\mathrm{\\mu G}$. These strong fields have also recently been inferred from observations \\citep{Longair:1994p901,Berezhko:2003p504,Vink:2003p515}; their existence, along with the requirement of stronger fields for the DSA mechanism, suggests that a magnetic field amplification mechanism is in operation. A possible amplification mechanism through a non-resonant instability was suggested in \\citet{Bell:2004p24}, following previous work on the resonant mode in \\citet{Lucek:2000p25}, and later extended to include multidimensional effects in \\citet{Bell:2005p27}. The non-resonant streaming instability, part of a class of streaming instabilities derived in \\citet{winskeleroy}, can be described by considering a MHD model for the background plasma, and a CR induced current imposed externally; the feedback of the electromagnetic fields on the CR particles is thus neglected. Although in the linear stage of the instability $\\lambda_\\mathrm{max}\\ll r_\\mathrm{LCR}$ ($\\lambda_\\mathrm{max}$ the fastest growing wavelength and $r_\\mathrm{LCR}$ the typical Larmor radius of the CRs), recent full PIC simulations by \\citet{Niemiec:2008p18}, \\citet{Riquelme:2009p190}, \\citet{reville} and \\citet{stroman} indicate that the feedback mechanism of the fields on the CRs is important in the nonlinear stage, and suggest that a careful study of the instability in this regime is important to determine the saturation levels of the magnetic field. Recent works by \\citet{Amato:2009p20} and \\citet{Luo:2009p19}, using kinetic theory, have also shown how the saturation of this non-resonant mode depends on the details of the particle distributions. The analysis of the behavior of the non-linear stage of the instability is very complex; full assessment of the saturation level of the magnetic field would imply a first-principle calculation of the shock formation, and the long term evolution of the shock precursor, including the self consistent acceleration of particles. The saturation of the instability was thus first assessed numerically in \\citet{Bell:2004p24}, where an external current was used to model the CRs, and an MHD model was used to simulate the background plasma. The saturation mechanism found was due to the tension of the magnetic field, which grows faster than the driving term; also, the size of the simulation box was seen to limit further growth of the instability \\citep{Bell:2004p24}. More recent PIC simulation results in \\citet{Niemiec:2008p18} actually show a saturation level of $\\delta B/B\\sim1$, but consider parameters such that $\\gamma_\\mathrm{max}/\\omega_\\mathrm{ci}\\ll1$ (the ratio of the growth rate of the fastest growing mode to the CR cyclotron frequency) is not strictly maintained, which is a requirement for the development of the non-resonant parallel mode. The full PIC simulation results in \\citet{Riquelme:2009p190} and \\citet{stroman}, taking into account the feedback of the electromagnetic fields on the CR particles, and using $m_i/m_e$ ratios up to $100$ and $n_\\mathrm{CR}/n_\\mathrm{b}$ down to $4\\times10^{-3}$, imply a saturation level of $\\delta B/B\\sim10$, which occurs when the relative drift between the CRs and the background plasma decreases. Also, in \\citet{reville}, similar saturation levels for the magnetic field are found. For PIC simulations, therefore, the instability saturates when CRs loose some of their bulk momentum to the background plasma, which starts to drift more rapidly in the direction of CRs until the velocities of the two species converge and the driver for the instability is eliminated. As shown in kinetic modeling by \\citet{Luo:2009p19}, the reduction of the CR streaming motion is due to CR resonant diffusion in non-resonant magnetic turbulence, which has been recently observed in \\citet{stroman}. The nonlinear evolution of the turbulence observed in fully kinetic simulations is also in qualitative agreement with the predictions of quasi-linear calculations for non-resonant modes derived for non-relativistic beams in \\citet{winskeleroy}. However, the accurate modeling of the late time evolution of the system, beyond the saturation, is limited in the kinetic simulations, when the assumption of plasma homogeneity is no longer valid. As argued in \\citep{reville}, when the background plasma velocity approaches the CR bulk velocity, the density in a real scenario also changes, modifying the underlying conditions for the simulation model. This reinforces the necessity for a thorough analysis of the behavior of this instability in its non-linear stage, covering a wide range of dynamical scales, and exploring the saturation limits in different regimes. Here, we present multi-dimensional simulation results of the non-resonant streaming instability, using the kinetic ion fluid electron hybrid model implemented in \\textit{dHybrid} (see \\citet{lgargate} for numerical implementation details). An important advantage of the hybrid simulations is to enable the study of the instability on the ion time scale, neglecting the high-frequency modes associated with the electrons; realistic density ratios down to $n_\\mathrm{CR}/n_\\mathrm{b}=10^{-5}$ can then be used, along with realistic ion-to-electron mass ratios, and simulations can be run to a point well beyond the linear stage, into the saturated state. The variation of the density ratio has an impact on the behavior of the non-resonant mode, with larger values implying the generation of different modes \\citep{Bell:2004p24,Niemiec:2008p18,Riquelme:2009p190}. Though the saturation mechanism is independent of the initial linear instability \\citep{stroman}, these modes lead to smaller saturation amplitudes and thus may not be directly relevant for the typical SNR shock scenarios. Also, in the hybrid simulation results shown, both the background and the CR ions are modeled kinetically, which enables the study of the feedback mechanism on the CR population. Results can then be easily compared with both the external current-driven MHD simulations, where feedback on the CR population is not modeled, and the fully kinetic simulations, which use small ion-to-electron mass ratios, and larger $n_\\mathrm{CR}/n_\\mathrm{b}$ ratios. This paper is organized as follows. In section 2, the parameters used in the simulations are discussed, and the results concerning the evolution of the linear and nonlinear stages of the instability are presented. The saturation mechanism is discussed in section 3, and compared with the most recent results in the literature. Finally, in the last section, we present the conclusions and outline future research directions. ", "conclusions": "The identification of the nonlinear saturation mechanism for the non-resonant streaming instability is important for the determination of the maximum magnetic field amplification in SNR shock scenarios. Amplification factors of $B_\\perp/B_\\parallel\\sim10$, with local temporary peaks of $B_\\perp/B_\\parallel\\sim25$ are observable in the current hybrid simulations; similar results were also recently shown in the kinetic simulations of \\citet{Riquelme:2009p190} \\citet{reville} and \\citet{stroman}, with reduced temporal and spatial scales, mass ratios and density ratios. The hybrid simulation results presented expand the current knowledge of the non-resonant streaming instability by extending the dynamical range under study, and by performing a detailed study of the instability under different driving regimes. Leveraging on the hybrid simulations presented here, it is possible to scan the parameter space and analyze the nonlinear stage of the instability in detail, using realistic $n_\\mathrm{CR}/n_\\mathrm{b}$ and $m/m_e$ ratios. Close examination of Fig. \\ref{fig3} and Fig. \\ref{fig4} shows a number of important points, relating to the saturation mechanism. As the $n_\\mathrm{CR}/n_\\mathrm{b}$ ratio is decreased towards $10^{-5}$, the energy in the CRs increases, as in our setup the current is initialized at the same level in all simulations; the increase in the free energy of the CRs, associated with the higher drift velocity $v_\\mathrm{sh}$, does not change the peak energy level of the magnetic field. The excess free energy is instead transferred to the background plasma, and also to the perpendicular velocity components of the CR population (Fig. \\ref{fig3} b), run $\\mathcal{C}_1$). The reaction of the background plasma is critical in the nonlinear stage of the instability. The average local Alfv\\'en velocity ($v_A\\propto B_\\parallel/\\sqrt{n}$ calculated in each cell and averaged over the entire simulation box) is approximately constant over the linear development of the instability. In the nonlinear stage, $v_A$ increases with the local magnetic field $B_\\parallel$, and with the background density variations (which are in phase), and thus $v_A\\propto\\sqrt{\\delta B/B_\\parallel}$. The background plasma accelerates in the CR propagation direction, and the current density peaks, marking the end of the linear stage (Fig. \\ref{fig4}). At this point in time, the CRs $v_x$ velocity component is still $v_x\\sim1500\\,\\mathrm{v_A}$, well above the Alfv\\'en velocity and above the background plasma drift velocity, which is sub-Alfv\\'enic. This shows that the saturation is not dependent on the amount of free energy in the CRs, but instead results from the reduction of the scale separation of the background plasma, and the CRs. The CR's $J_\\perp$ increases for run $\\mathcal{B}_2$, Fig. \\ref{fig4} d), as it is substantially lower than $J_x$ initially, and the distribution becomes isotropic after the saturation, at $t\\sim15\\,\\mathrm{\\gamma^{-1}}$. For run $\\mathcal{B}_1$ in Fig. \\ref{fig4} c), the distribution is also isotropic at the end, but the CRs $J_\\perp$ velocity component is not affected. Increasing the free energy in the CRs does not affect the saturation level of the magnetic field significantly, as long as the driving current is maintained. The hybrid simulation results presented thus show that, for a broad range of parameters, the nonlinear growth of the streaming instability is independent on the energy of the driving CR population, and depends only on the current carried by the CRs. The amplification factor of $\\sim10$ for the perpendicular magnetic field above the seed $B_\\parallel$ feed indicates that the mechanism is relevant for the magnetic field amplification by CR particles in SNR shocks. Beyond the saturation point, for $t>16\\,\\mathrm{\\gamma^{-1}}$, the CR particle distribution becomes isotropic, with the instability saturating after 10 e-foldings. The de-magnetization of the background plasma, as it gains energy, hinders further growth of the non-resonant mode; when both the CRs and the background plasma are unmagnetized, the instability should behave more like the Weibel instability \\citep{weibel}. Our results thus indicate that the instability begins to saturate when the current carried by the background ions is similar to the current carried by the CR population (see Fig. \\ref{fig4}). This results in a final average velocity for the background plasma of $v_\\mathrm{b}\\sim n_\\mathrm{CR}/n_\\mathrm{b} v_\\mathrm{sh}$ so that $V_\\mathrm{b}\\sim10^{-5}v_\\mathrm{sh}$ for realistic parameters. At this point, the instability is in the non-linear stage, and the magnetic field growth rate is much lower ($t>9\\,\\mathrm{\\gamma^{-1}}$ in Fig. \\ref{fig3} c) and Fig \\ref{fig4} a) and c) for run $\\mathcal{B}_1$). In fact, even in the linear stage of the instability, energy is being transferred to the perpendicular components of the background plasma at a greater rate than into the magnetic field perpendicular components (see Fig. \\ref{fig3} c), and thus at some point the background plasma de-magnetizes and the instability saturates. Previous simulation PIC results \\citep{Riquelme:2009p190,reville,stroman} show similar saturation levels for the magnetic field $\\delta B/B\\sim10$: the instability saturates when the background plasma is moving with a bulk velocity comparable to the CR population drift speed, which is equivalent to saying that the currents carried by the background plasma and the CRs are similar (since $n_\\mathrm{CR}\\sim n_\\mathrm{b}$, and thus $J_\\mathrm{CR}=n_\\mathrm{CR}\\,e\\,v_\\mathrm{sh}\\sim J_\\mathrm{b}=n_\\mathrm{b}\\,e\\,v_\\mathrm{b}$). It has been claimed \\citep{reville} that this saturation limit might be numerical rather than physical because in a real shock precursor a significant variation in the background plasma velocity should be concurrent with a variation in the background plasma density.However, here we show for $m_i\\gg m_e$ and $n_\\mathrm{CR}\\ll n_\\mathrm{b}$ that the instability saturates due to the de-magnetization of the background plasma; at that point the currents carried by the background ions and the CRs are comparable, although the velocities differ significantly (since in our hybrid simulations $n_\\mathrm{CR}\\ll n_\\mathrm{b}$ implies that $v_\\mathrm{b}\\sim v_\\mathrm{sh}\\,n_\\mathrm{CR}/n_\\mathrm{b}$ from the equality of the currents at the saturation). Our simulations thus reinforce the result $\\delta B/B\\sim10$ at saturation, and indicate that the instability should be important in a young SNR shock precursor. The simulations present in the literature make assumptions about the physical scenarios, and therefore significant care should be taken when interpreting and comparing the results. For instance, the back-reaction of the fields in the CRs is not accounted for when using external currents as drivers, in full PIC simulations mass ratios and density ratios far from the physical conditions have been explored, and the dynamics of electrons is not accounted for in hybrid simulations. Thus, our results further indicate that first-principle simulations of SNR shocks and particle acceleration should be attempted in a future work, since the different simulation methods can capture the instability, but differences might still be found when the simulation setup and the models are refined in order to more closely match the experimental conditions. Further analysis for SNR shocks will be possible by considering the saturation mechanism when unperturbed CRs are continuously injected into the simulation box. This will allow for an improved comparison of the SNR shock scenario, where CR particles are also thought to be continuously injected from the shock into the upstream medium. This will be explored in a future publication, leveraging on the unique characteristics of the hybrid model. Finally, with hybrid simulations, it will also be possible to determine the dependence of usual characteristics of the instability with the distance to the shock, and to study in detail the energy profile of the accelerated CR particles." }, "1002/1002.1647_arXiv.txt": { "abstract": "The type Ib supernova 2010O was recently discovered in the interacting starburst galaxy Arp 299. We present an analysis of two archival \\emph{Chandra} X-ray observations of Arp 299, taken before the explosion and show that there is a transient X-ray source at a position consistent with the supernova. Due to the diffuse emission, the background is difficult to estimate. We estimate the flux of the transient from the difference of the two X-ray images and conclude that the transient can be described by a \\rev{0.225 keV black body with a luminosity of $2.5\\pm0.7 \\times 10^{39}$ erg s$^{-1}$} for a distance of 41 Mpc. These properties put the transient in between the Galactic black hole binary XTE J1550-564 and the ultra-luminous X-ray binaries NGC 1313 X-1 and X-2. The high level of X-ray variability associated with the active starburst makes it impossible to rule out a chance alignment. If the source is associated with the supernova, it suggests SN2010O is the explosion of the second star in a Wolf-Rayet X-ray binary, such as Cyg X-3, IC 10 X-1 and NGC 300 X-1. ", "introduction": "Type Ib supernovae lack hydrogen in their spectrum, but show helium lines, while Ic supernovae also lack helium. They are regarded as the core-collapse explosions of massive stars that have lost their hydrogen envelopes and thus have become Wolf-Rayet stars \\citep{1986ApJ...306L..77G}. These Wolf-Rayet stars can originate from the most massive stars that can remove their hydrogen envelope on and just after the main sequence in strong stellar winds, or from lower-mass stars via a binary interaction that removes the hydrogen envelope \\citep[e.g.][]{1992ApJ...391..246P,1995PhR...256..173N}. In the last decade a new way of linking supernovae with their progenitors has become available. The growing archive of high-resolution images has made it possible to detect the progenitors of type II (hydrogen rich) supernovae in pre-explosion optical images \\citep[see][for a review]{2009ARA&A..47...63S}. No progenitors of type Ib or Ic supernovae have been found, despite deep pre-explosion images for 10 of them. However, the relative frequency of Ib/Ic to type II supernovae of around 0.4 suggests that binary interactions play an important role, as for a standard initial mass function, there are not enough very massive stars that could lose their envelope via a stellar wind \\citep{2009ARA&A..47...63S}. An alternative method for directly detecting supernova progenitors is to use archival images in other wave bands, such as X-ray data. We started a program to search for type Ia supernova progenitors in \\emph{Chandra} X-ray data and have found one likely progenitor and derived five upper limits \\citep{vn08,rbv+08,nvr+08}. For type Ia supernova progenitors the argument to look for X-ray progenitors is the suggestion that supersoft X-ray sources may produce type Ia supernovae, based on the accreting white dwarf model \\citep{1973ApJ...186.1007W,nom82}. For type Ib and Ic supernovae, the binary progenitor scenario suggests the possibility that the Ib or Ic explosion is the second supernova in the binary and thus before the explosion may have been part of a high-mass X-ray binary. Indeed, the relative frequency estimates of \\citet[][their figure 16]{1992ApJ...391..246P} suggest that the likelihood for a Ib to be the second explosion in the binary is at least as high as it being the first one. It is therefore useful to constrain the X-ray properties of the direct progenitors of Ib and Ic supernovae. Even more tantalising evidence for such scenario comes from the known Wolf-Rayet star in the X-ray binary Cyg X-3 \\citep{1996A&A...314..521V} and the recent discovery of two X-ray binaries in which a (massive) Wolf-Rayet star orbits a black hole, IC 10 X-1 and NGC 300 X-1 \\citep{2004A&A...414L..45C,2004ApJ...601L..67B,2007ApJ...669L..21P,2007A&A...461L...9C}. Given the short life times of massive Wolf-Rayet stars, a Ib or Ic supernova explosion is expected within next few million years in these systems. In this paper we report the discovery of a transient soft X-ray source at the position of the recent Ib supernovae 2010O that exploded in the interacting starburst galaxy Arp 299 (\\rev{IC 694}/NGC 3690). In Sect.~\\ref{2010O} we discuss supernova 2010O and its host galaxy, in Sect.~\\ref{obs} the X-ray, optical and radio observations that we used and in Sect.~\\ref{results} the results of the \\emph{Chandra} analysis. In Sect.~\\ref{discussion} we discuss the implications of the finding, the possibility of a chance alignment and the scope for future work. ", "conclusions": "\\label{discussion} If the X-ray transient found in the second \\emph{Chandra} image is associated with the SN2010O (see below for a discussion of the possibility of a chance alignment), it suggests that this is the explosion of the second star in a binary system. Before the supernova explosion the system must have consisted of a Wolf-Rayet star orbiting a compact object. We can compare the X-ray properties of the transients to the outbursts seen in Galactic black-hole X-ray binaries \\citep[see][for an overview]{mr04}. The peak luminosity is higher than those of the Galactic sources, but these are particularly poorly known, as the distances to these systems are not well determined \\citep[see][]{jn04}. The outburst luminosity is below the Eddington limit of a 10 \\msun black hole, if we assume pure helium accretion. The spectrum is much softer than the typical 1 keV found in Galactic black hole binaries. However, the very soft spectrum and high luminosity fit very well in between the highest luminosity point of the Galactic X-ray binary XTE J1550-564 and the ultra-luminous X-ray binaries NGC 1313 X-1 and X-2 \\citep[see figure 7 of][]{2007Ap&SS.311..213S}. The transient thus could be a relative of the nearby Wolf-Rayet X-ray binaries Cyg X-3 \\citep{1996A&A...314..521V}, IC 10 X-1 \\citep{2004A&A...414L..45C,2004ApJ...601L..67B} and NGC 300 X-1 \\citep{2007A&A...461L...9C} and the soft spectrum suggests a (massive) black hole accretor. However, all of these seem to be fairly persistent X-ray sources, with luminosities well below the Eddington limit. The transient thus would represent the equivalent of the outburst in soft X-ray transients, in a Wolf-Rayet X-ray binary. With a distance modulus of 33 and and a reported visual magnitude of 15.8 around the peak \\citep{2010CBET.2144....2N}, together with the extinction of 0.65 as derived from the spectrum \\citep{2010CBET.2149....1M}, the absolute magnitude of SN2010O is $-17.9$. This is well within the large range of absolute magnitudes of type Ib supernovae \\citep{2006AJ....131.2233R}. Unfortunately, there are no known correlations of Ib supernova explosion features and the properties of the progenitor Wolf-Rayet stars. The whole discussion above depends on the association of the X-ray transient with the supernova. With WAVDETECT 22 unique sources are found within half an arc min from the position of the supernova. If we assume we would consider sources within 1$\\arcsec$ as a possible match that would give a chance alignment probability of 2 per cent. Ten of these sources are actually within 15$\\arcsec$ and the difference image in Fig.~\\ref{fig:diff} shows \\rev{two sources not detected by WAVDETECT if we include the transient}. We therefore conclude that there is a 5 per cent probability of a chance alignment. However, the close match in position and general consistency of the X-ray properties with what could be expected from an X-ray binary in which the second star exploded as a Ib supernova, make the connection plausible. \\rev{The possible association of the supernova with the cluster seen in the \\emph{HST} ACS observations \\citep{2010ATel.2422....1B} also allows for the possibility that the X-ray source is coincident with the position of the SN, but is actually an unrelated X-ray binary in the cluster. However this is not very likely, as there are only 18 X-ray sources found in Arp 299, and many more star clusters, and a study by \\citet{2004MNRAS.348L..28K} shows that in starburst galaxies the clusters have only rarely an X-ray source right on top of them}. We therefore conclude that the observations presented in this paper suggest the exciting possibility that supernova 2010O was the explosion of a massive Wolf-Rayet star that was part of a X-ray binary containing a black hole. Further study, in particular for other supernovae with pre-supernova \\emph{Chandra} data, will be needed to address the question of how often type Ib and Ic supernovae explode in X-ray binaries. We are in the process of analysing the earlier type Ib and Ic supernovae for which there is archival \\emph{Chandra} data (Nielsen et al. in preparation)." }, "1002/1002.3368_arXiv.txt": { "abstract": "The first stars in the history of the Universe are likely to form in the dense central regions of $\\sim 10^5$--$10^6\\ M_\\odot$ cold dark matter halos at $z\\approx 10$--50. The annihilation of dark matter particles in these environments may lead to the formation of so-called dark stars, which are predicted to be cooler, larger, more massive and potentially more long-lived than conventional population III stars. Here, we investigate the prospects of detecting high-redshift dark stars with the upcoming James Webb Space Telescope (JWST). We find that all dark stars with masses up to $10^3\\ M_\\odot$ are intrinsically too faint to be detected by JWST at $z>6$. However, by exploiting foreground galaxy clusters as gravitational telescopes, certain varieties of cool ($T_\\mathrm{eff}\\leq 30000$ K) dark stars should be within reach at redshifts up to $z\\approx 10$. If the lifetimes of dark stars are sufficiently long, many such objects may also congregate inside the first galaxies. We demonstrate that this could give rise to peculiar features in the integrated spectra of galaxies at high redshifts, provided that dark stars make up at least $\\sim 1\\%$ of the total stellar mass in such objects. ", "introduction": "\\label{intro} The first stars in the history of the Universe are predicted to form inside $\\sim$10$^5$--$10^6 M_\\odot$ minihalos at redshifts $z\\approx 10$--50 \\citep[e.g.][]{Tegmark et al.,Yoshida et al.}. Due to the lack of efficient coolants in the primordial gas at these early epochs, the resulting population III stars are believed to be very massive \\citep[$\\gtrsim 100\\ M_\\odot$; e.g.][]{Abel et al.,Bromm et al. b}, hot \\citep[effective temperature $T_\\mathrm{eff}\\sim 10^5$ K; e.g][]{Bromm et al. a} and short-lived \\citep[$\\approx 2$--3 Myr;][]{Schaerer a}. Non-rotating population III stars with masses of 50--140 $M_\\odot$ or $M>260 \\ M_\\odot$ are expected to collapse directly to black holes, whereas stars with masses of 140--$260 \\ M_\\odot$ may produce luminous pair-instability supernovae \\citep[e.g.][]{Heger et al.}. The latter may enrich the ambient medium with heavy elements and initiate the transition to the normal mode of star formation (population I and II, with a characteristic stellar mass $<1 \\ M_\\odot$) known from the low-redshift Universe. The highly energetic radiation emitted from population III stars during their lifetimes may also have played an important role in cosmic reionization at $z>6$ \\citep[e.g.][]{Sokasian et al.,Trenti & Stiavelli}. An observational confirmation of very massive population III stars would be an important breakthrough in the study of the star formation, chemical enrichment and reionization history of the Universe. The James Webb Space Telescope\\footnote{http://www.jwst.nasa.gov}(JWST), scheduled for launch in 2014, has been designed to study the epoch of the first light, reionization and galaxy assembly, but is not expected to be able to directly detect individual population III stars at $z\\gtrsim 10$. Searches for population III stars with the JWST would instead focus on the pair-instability supernovae produced at the end of their lifetimes \\citep{Weinmann & Lilly}, or on $\\sim$10$^5$--$10^7\\ M_\\odot$ clusters of population III stars \\citep{Bromm et al. a,Scannapieco et al.,Trenti et al.,Johnson et al. b,Johnson}. Other alternatives are to look for the spectral signatures of population III stars forming in pockets of unenriched gas within high-redshift galaxies \\citep{Tumlinson & Shull,Schaerer a,Schaerer b,Dijkstra & Wyithe,Johnson et al. a}, or their integrated contribution to the infrared extragalactic background light \\citep{Santos et al.,Cooray et al.}. It has recently been recognized that annihilation of dark matter in the form of Weakly Interacting Massive Particles (WIMPs; e.g.~the lightest supersymmetric or Kaluza-Klein particles, or an extra inert Higgs boson) may have generated a first population of stars with properties very different from the canonical population III \\citep{Spolyar et al. a}. Because the first stars are likely to form in the high-density central regions of minihalos, annihilation of dark matter into standard model particles could serve as an additional energy source alongside or instead of nuclear fusion within these objects. This leads to the formation of so-called dark stars, which are predicted to be cooler, larger, more massive and potentially longer-lived than conventional population III stars \\citep{Spolyar et al. a,Iocco b,Freese et al. a,Iocco et al.,Yoon et al.,Taoso et al.,Natarajan et al.,Freese et al. b,Spolyar et al. b,Umeda et al.,Ripamonti et al. a}. Similar effects have been seen in studies of the impacts of dark matter upon population I and II stars \\citep{Salati & Silk,Fairbairn et al.,Scott et al. a,Scott et al. b,Casanellas & Lopes}. A significant population of high-redshift dark stars could have important consequences for the formation of intermediate and supermassive black holes \\citep{Spolyar et al. b}, for the cosmic evolution of the pair-instability supernova rate \\citep{Iocco c}, for the X-ray extragalactic background and for the reionization history of the Universe \\citep{Schleicher et al.}. Effects such as these can be used to indirectly constrain the properties of dark stars, but no compelling evidence for or against a dark star population at high redshifts has so far emerged. Here, we explore a more direct approach -- the prospects for detection of population III dark stars using the JWST. When attempting to assess the detectability of dark stars at high redshifts, the expected lifespan of such objects represents a crucial aspect. In principle, dark stars could live indefinitely, provided that there is ample dark matter available to fuel them. Dark stars are powered by gravitationally-contracted dark matter, pulled into their core as infalling gas steepens the gravitational potential during the formation phase. Because annihilation depletes the dark matter present within a star, ordinary fusion processes eventually take over as the dominant power source if the dark matter is not replenished. At this point, the dark star will essentially transform into a conventional population III star, albeit more massive because the increased duration of the formation phase has allowed it to accrete more gas. The dark matter present within the star during formation will last only for a few million years, but these reserves may later be replenished by scattering of WIMPs on nucleons, causing them to lose energy and become captured in the stellar core. This could boost the longevity of dark stars substantially \\citep{Iocco et al.,Freese et al. a-1,Yoon et al.,Taoso et al.,Spolyar et al. b,Iocco c}. The ongoing replenishment of the dark matter through capture and the resultant increase in longevity rely upon a number of strong assumptions and approximations, and the feasibility of such a mechanism is still yet to be proven by detailed calculations \\citep{Sivertsson & Gondolo}. Whether the capture process will be efficient depends on two factors: the scattering cross-sections of the dark matter particles with ordinary nucleons, and the amount and density of dark matter available for capture from the star's surroundings. The WIMP-nucleon scattering cross-sections can be constrained by direct detection experiments \\citep[e.g.][]{Savage et al.,CDMS08} and searches for neutrinos produced by annihilation in the Sun \\citep[e.g.][]{IceCube09}. However, the question of the amount of dark matter available for refueling is more complicated \\citep{Sivertsson & Gondolo}. In a pristine halo, WIMP annihilations and scatterings during the formation stage would eventually deplete orbits with low angular momenta and result in a cavity of reduced dark matter density in the vicinity of the dark star. Whether infalling WIMPs could restore the balance before the star evolves into a supernova or a black hole may depend on the overall structural evolution of the minihalo (e.g. contraction, mergers with other halos), on tidal interactions of the dark star with subhalos, gas clouds and possibly other population III stars within the minihalo itself. Should a violent event cause the dark star to venture far from the centre of the minihalo, the dark matter density would very quickly become too low to sustain further dark matter burning. At the current time, estimates of the dark star lifetime in the presence of capture range from a few times $10^5$ to $10^{10}$ years \\citep[e.g.][]{Yoon et al.,Iocco c,Sivertsson & Gondolo}. Many of the mechanisms mentioned above for replenishing the dark matter in the centre of the dark star have moreover not yet been explored. In this paper, we will therefore treat the duration of the dark star phase as a free parameter. Model atmospheres and evolutionary histories of dark stars are presented in Sect.~\\ref{models}. The detectability of isolated dark stars, with and without the effects of gravitational lensing by foreground galaxy clusters, is explored in Sect.~\\ref{detectability}. In Sect.~\\ref{discussion}, we explain how high-redshift dark stars can be distinguished from other objects based on their JWST colours, and discuss the possibility of detecting the spectral signatures of dark stars in the first generations of galaxies. Sect.~\\ref{summary} summarizes our findings. Throughout this paper, we will assume a $\\Lambda$CDM cosmology with $\\Omega_\\Lambda=0.73$, $\\Omega_M=0.27$ and $H_0=72$ km s$^{-1}$ Mpc$^{-1}$. After this paper had been submitted, another paper \\citep{Freese et al. c} dealing with the prospects of detecting high-redshift dark stars with the JWST was posted on the arXiv preprint server. The two studies differ in their assumptions concerning the masses of dark stars. Whereas we consider objects with masses up to $\\sim 10^3\\ M_\\odot$, Freese et al. instead focus on the detectability of `supermassive dark stars' (with masses up to $10^7\\ M_\\odot$). \\begin{deluxetable*}{llllllllll} \\tabletypesize{\\scriptsize} \\centering \\tablecaption{Dark star models\\label{modeltable}} \\tablewidth{0pt} \\tablehead{ \\colhead{WIMP mass} & \\colhead{$M$ (M$_\\odot$)\\tablenotemark{a}} & \\colhead{$R$ (cm)\\tablenotemark{a}} & \\colhead{$\\log_{10}(g)$\\tablenotemark{b}} & \\colhead{$T_\\mathrm{eff}$ (K)\\tablenotemark{a}} & \\colhead{$t_\\mathrm{burn}$ (yr)} & \\colhead{$t_\\mathrm{SA}$ (yr)} & \\colhead{$t_\\mathrm{max}$ (yr)} & \\colhead{Atmosphere} & \\colhead{$\\max(z_\\mathrm{obs})$\\tablenotemark{c}} } \\startdata 1\\, GeV &106 & $2.4 \\times 10^{14}$ & $-$0.612 & $5.4 \\times 10^3$ & $>10^{10}$ & $2.0 \\times 10^{5}$ & $2.0 \\times 10^{7}$ & \\textsc{marcs} & 10 \\\\ &371 & $2.7 \\times 10^{14}$ & $-$0.170 & $5.9 \\times 10^3$ & $>10^{10}$ & $1.2\\times 10^{6}$ & $1.2\\times 10^{8}$ & \\textsc{marcs} & 11 \\\\ &690 & $1.1 \\times 10^{14}$ &\\phantom{$-$}0.879 & $7.5 \\times 10^3$ & $>10^{10}$ &$2.4\\times 10^{7}$ & $5.0\\times10^8$ & \\textsc{marcs} & 11 \\\\ &756 & $3.7 \\times 10^{13}$ &\\phantom{$-$}1.865 & $1.0 \\times 10^4$ & $>10^{10}$ & $2.3 \\times 10^{8}$ & $5.0\\times10^8$ & \\textsc{tlusty} & 9 \\\\ &793 & $5.7 \\times 10^{12}$ &\\phantom{$-$}3.511 & $3.0 \\times 10^4$ & $>10^{10}$ & $8.7\\times 10^{8}$ & $5.0\\times10^8$ & \\textsc{tlusty} & 11\\\\ &824 & $5.8 \\times 10^{11}$ &\\phantom{$-$}5.512 & $1.1 \\times 10^5$ & $6.2 \\times 10^{6}$ & $4.9\\times 10^{9}$ & $6.2 \\times 10^{6}$ & \\textsc{tlusty} & 0.5\\\\ 100\\,GeV &106 & $7.0 \\times 10^{13}$ &\\phantom{$-$}0.458 & $5.8 \\times 10^3$ & $>10^{10}$ & $6.0\\times 10^{6}$ & $5.0\\times10^8$ & \\textsc{marcs} & 6\\\\ &479 & $8.4 \\times 10^{13}$ &\\phantom{$-$}0.955 & $7.8 \\times 10^3$ & $>10^{10}$ & $2.2\\times 10^{7}$ & $5.0\\times10^8$ & \\textsc{marcs} & 11\\\\ &716 & $1.1 \\times 10^{13}$ &\\phantom{$-$}2.895 & $2.3 \\times 10^4$ & $>10^{10}$ & $2.9\\times 10^{8}$ & $5.0\\times10^8$ & \\textsc{tlusty} & 13\\\\ &756 & $2.0 \\times 10^{12}$ &\\phantom{$-$}4.399 & $5.5 \\times 10^4$ & $2.2 \\times 10^{9}$ & $1.6\\times 10^{9}$ & $5.0\\times10^8$ & \\textsc{tlusty} & 3\\\\ &787 & $5.8 \\times 10^{11}$ &\\phantom{$-$}5.492 & $1.1 \\times 10^5$ & $5.9 \\times 10^{6}$ & $4.5\\times 10^{9}$ & $5.9 \\times 10^{6}$ & \\textsc{tlusty} & 0.5\\\\ 10\\,TeV &106 & $2.2 \\times 10^{13}$ &\\phantom{$-$}1.463 & $6.0 \\times 10^3$ & $>10^{10}$ & $1.7\\times 10^{8}$ & $5.0\\times10^8$ & \\textsc{marcs} & 2\\\\ &256 & $2.2 \\times 10^{13}$ &\\phantom{$-$}1.846 & $8.0 \\times 10^3$ & $>10^{10}$ & $3.1\\times 10^{8}$ & $5.0\\times10^8$ & \\textsc{marcs} & 4\\\\ &327 & $2.0 \\times 10^{13}$ &\\phantom{$-$}2.036 & $1.0 \\times 10^4$ & $>10^{10}$ & $2.8\\times 10^{8}$ & $5.0\\times10^8$ & \\textsc{tlusty} & 5\\\\ &399 & $6.6 \\times 10^{12}$ &\\phantom{$-$}3.085 & $2.5 \\times 10^4$ & $>10^{10}$ & $2.9\\times 10^{8}$ & $5.0\\times10^8$ & \\textsc{tlusty} & 7\\\\ &479 & $2.9 \\times 10^{12}$ &\\phantom{$-$}3.879 & $3.2 \\times 10^4$ & $>10^{10}$ & $1.9\\times 10^{9}$ & $5.0\\times10^8$ & \\textsc{tlusty} & 2\\\\ &550 & $6.0 \\times 10^{11}$ &\\phantom{$-$}5.307 & $9.5 \\times 10^4$ & $1.3 \\times 10^{7}$ & $3.6\\times 10^{9}$ & $1.3 \\times 10^{7}$ & \\textsc{tlusty} & 0.5\\\\ &553 & $4.8 \\times 10^{11}$ &\\phantom{$-$}5.503 & $1.1 \\times 10^5$ & $5.1 \\times 10^{6}$ & $3.9\\times 10^{9}$ & $5.1 \\times 10^{6}$ & \\textsc{tlusty} & $<0.5$\\\\ \\enddata \\tablenotetext{a}{from \\protect\\citet{Spolyar et al. b}} \\tablenotetext{b}{in units of g cm$^{-1}$ $s^{-2}$} \\tablenotetext{c}{This observability limit is set by the requirement that a single dark star should be sufficiently bright in at least one JWST filter to give a $5\\sigma$ detection after a 100 h exposure, if a gravitational magnification of $\\mu=160$ is assumed (see Sect.~\\ref{detectability}).} \\end{deluxetable*} ", "conclusions": "\\label{discussion} \\begin{figure} \\plotone{f5.eps} \\caption{The JWST/NIRCam $m_{356}-m_{444}$ vs. $m_{200}-m_{277}$ colours of $T_\\mathrm{eff}<10000$ K dark stars at $z=10$ (red star symbols) compared to a number of potential interlopers in multiband surveys: star clusters or galaxies at $z=0$--15 (black dots), AGN template spectra at $z=0$--15 (yellow dots), Milky Way stars with $T_\\mathrm{eff}=2000$-50000 K and $Z=0.001-0.020$ (blue dots) and Milky Way brown dwarfs with $T_\\mathrm{eff}=130$--2200 K (green dots). Since the dark stars reside a region of this colour-colour diagram that is disconnected from those occupied by these other objects, it should be possible to identify possible dark star candidates in deep multiband JWST/NIRCam surveys. \\label{colour_criteria}} \\end{figure} \\subsection{How to distinguish isolated dark stars from other objects} As demonstrated in Sect.~\\ref{lensing}, certain varieties of $z\\approx 10$ dark stars may be sufficiently bright and numerous to be detected by a JWST/NIRCam survey of the high-magnification regions of a foreground galaxy cluster. But how does one identify such objects among the overwhelming number of mundane interlopers located in front of, inside or beyond the lensing cluster? Given the many degeneracies involved in the interpretation of broadband photometry, there may well be unresolvable ambiguities in some cases. However, many of the cooler high-redshift dark stars should stand out in multiband survey data because of their unusual colours. This is demonstrated in Fig.~\\ref{colour_criteria}, where we plot the colour indices $m_{356}-m_{444}$ vs. $m_{200}-m_{277}$ (based on AB-magnitudes in the JWST/NIRCam F200W, F277W, F356W and F444W filters) at $z=10$ for all dark stars from Table~\\ref{modeltable} with $T_\\mathrm{eff}< 10000$ K. The colours of these models (red star symbols) are compared to the colours predicted for a wide range of galaxies, star clusters, active galactic nuclei (AGN) and Milky Way stars. The cloud of black dots in Fig.~\\ref{colour_criteria} indicate the colours of integrated stellar populations (star clusters and galaxies) generated with the \\citet{Zackrisson et al.} spectral synthesis model. These predictions are based on instantaneous-burst\\footnote{This is a conservative choice, since allowing for more extended star formation histories would only result in a more restricted colour coverage for these objects}, Salpeter-IMF stellar populations at redshifts $z=0$--15 with metallicities in the range $Z=0.001$-0.020, ages ranging from $10^6$ yr up to the age of the Universe at each redshift and a rest-frame stellar dust reddening of $E(B-V)=0$--0.5 mag assuming the \\citet{Calzetti et al.} extinction law. Also included in Fig.~\\ref{colour_criteria} are the expected colours of foreground stars with $T_\\mathrm{eff}=2000$-50000 K and $Z=0.001-0.020$ in the Milky Way (i.e. at $z=0$), based on the \\citet{Lejeune et al.} compilation of synthetic stellar atmosphere spectra (blue dots), and the colours of Milky Way brown dwarfs in the 130--2200 K range based on the \\citet{Burrows et al. a,Burrows et al. b} models (green dots). The yellow dots represent the template AGN spectra of \\citet{Hopkins et al.} for bolometric luminosities $\\log_{10} L_\\mathrm{bol}/L_\\odot=8.5$--14.0, at redshifts $z=0$--15. Since none of these potential interlopers have $m_{356}-m_{444}$ and $m_{200}-m_{277}$ colours that overlap with those of cool $z\\approx 10$ dark stars, a diagnostic diagram of this type can be used to cull objects that are clearly {\\it not} dark stars from multiband survey data. However, given that $T_\\mathrm{eff}< 10000$ K dark stars are unlikely to attain apparent magnitudes brighter than $m_{AB}\\approx 30$ at their peak wavelengths (even when boosted by gravitational lensing; see Fig.~\\ref{ABmag_nolens2}b), and can realistically only be detected in one or two NIRCam filters, follow-up spectroscopy of the remaining dark star candidates will be required to establish their exact nature. This will admittedly be very challenging, but a very coarse spectrum can possibly be obtained for $m_{AB}\\approx 30$ objects with JWST/NIRSpec (assuming a $3.6\\times 10^5$ s exposure). Objects of this type may also be suitable targets for the 42 m European Extremely Large Telescope\\footnote{http://www.eso.org/sci/facilities/eelt/} (E-ELT). \\subsection{The spectral signature of dark stars in high-redshift galaxies} The first galaxies are expected to form at $z\\approx 10$--13 in CDM halos with total masses around $10^8\\ M_\\odot$ \\citep[e.g.][]{Johnson et al. a,Johnson et al. b,Greif et al.,Ricotti}. At the time of assembly, each such object is likely to contain a number of minihalos in which population III stars have already formed \\citep{Greif et al.}. If some of these population III stars go through a long-lived ($\\tau\\gtrsim 10^8$ yr) dark star phase, several dark stars may in principle congregate inside the first generation of galaxies and give rise to telltale signatures in their integrated spectra. In Fig.~\\ref{firstgal}, we display the rest-frame spectrum predicted by the \\citet{Zackrisson et al.} population synthesis model for a $10^8$ yr old, low-metallicity ($Z=0.001$, i.e. population II), Salpeter-IMF (mass range 0.08--120 $M_\\odot$) stellar population which has formed stars at a constant rate. This population has been assigned a stellar mass of $10^6 \\ M_\\odot$, which -- for a $10^8\\ M_\\odot$ halo with baryon fraction $f_\\mathrm{bar}\\approx 0.7$ -- corresponds to $\\sim 10\\%$ of the baryonic mass in stars, in rough agreement with the models of \\citet{Ricotti} for a $z\\approx 10$ galaxy. The predicted NIRCam magnitudes of this object lie around $m_{AB}\\approx 33$--34 which would make it sufficiently bright for detection with JWST if seen through a gravitational lens with magnification $\\mu\\gtrsim 10$. Superposed on the spectrum of this high-redshift galaxy is the integrated contribution from ten $T_\\mathrm{eff}=7500$ K (the 690 $M_\\odot$ models from the 1 GeV WIMP track) dark stars (red line). These dark stars, which contribute only $0.7\\%$ of the stellar mass in this galaxy, give rise to a conspicuous red bump in the spectrum at rest-frame wavelengths longward of 0.36$\\mu$m (this corresponds to wavelengths longer than 3.96$\\mu$m at $z=10$). Because of this, galaxies that contain many cool dark stars are expected to display anomalously red colours. A feature like this is very difficult to produce through other means. For instance, the spectrum depicted in Fig.~\\ref{firstgal} cannot be attributed to dust reddening, since no known extinction law would allow a sharp rise in flux at wavelengths longward of 0.36$\\mu$m to co-exist with a very blue continuum at shorter wavelengths. Attributing the red bump as due to thermal dust emission is also untenable, since this would require dust radiating at a temperature close to that of the dark star (7500 K), which is higher than the sublimation temperature of all known types of dust. \\begin{figure} \\plotone{f6.eps} \\caption{The rest-frame spectrum of a $10^6\\ M_\\odot$, $Z=0.001$, Salpeter-IMF stellar population (stellar mass range 0.08--120 $M_\\odot$) which has formed stars at a constant rate for $10^8$ yr (black line) with a superimposed contribution (red line) from ten $T_\\mathrm{eff}=7500$ K, 690 $M_\\odot$ dark stars (from the 1 GeV WIMP track). Despite making up only $0.7\\%$ of the stellar mass in this galaxy, these dark stars give rise to an upturn in the spectrum at rest-frame wavelengths longward of 0.36$\\mu$m. In a $z=10$ object, this spectral feature should appear longward of 3.96$\\mu$m and could serve as a telltale signature of cool dark stars within the first galaxies. Due to the low but non-zero metallicity ($\\left[ \\mathrm{Fe} / \\mathrm{H} \\right] = -5$) of the MARCS model for the 690 $M_\\odot$ dark stars, numerous metal absorption lines are seen throughout the red spectrum. While these lines appear to be strong due to the high spectral resolution of the model, they actually have very small equivalent widths and negligible impact on the JWST broadband fluxes. \\label{firstgal}} \\end{figure} In Fig.~\\ref{colour_criteria2}, we display the $m_{356}-m_{444}$ vs. $m_{200}-m_{277}$ colour evolution (black solid line) as a function of age for the synthetic galaxy from Fig.~\\ref{firstgal} at $z=10$. Here, the age runs from $10^6$ yr (black triangle) up to the age of the Universe at this redshift ($\\approx 5\\times 10^8$ yr). Also indicated are the colours of two cool dark stars: the 7500 K model from the 1 GeV WIMP track (blue star) and the 5800 K model from the 100 GeV WIMP track (red star). Due to their low temperatures, these two dark stars are far redder than the model galaxy in both colours plotted. As shown in Fig.~\\ref{colour_criteria}, this is generally the case for $T_\\mathrm{eff}<10000$ K dark stars observed in these filters. The galaxy is here assumed to go through a short burst of star formation (forming stars at a constant rate for $10^8$ yr), after which it evolves passively. This gives a conservative estimate of the colour difference between galaxies and dark stars. Allowing a more extended star formation episode or a star formation rate that increases over time would only increase the discrepancy between the colours of dark stars and the model galaxy. The dashed lines indicate how the colours of a $10^8$ yr galaxy would shift in this diagram, if it were to harbour dark stars of the type considered. The filled circles along each such mixing track indicate the position at which the dark stars make up $1\\%$ of the total stellar mass. Since these points are significantly redder in the $m_{356}-m_{444}$ colour than the reddest point along the standard galaxy track, a $\\sim 1\\%$ stellar mass fraction in dark stars would result in very peculiar colours for $z=10$ galaxies and should allow such objects to be identified as candidate `dark star galaxies' in JWST multiband survey data. \\begin{figure} \\plotone{f7.eps} \\caption{The JWST/NIRCam $m_{356}-m_{444}$ vs. $m_{200}-m_{277}$ colour evolution as a function of age for a $z=10$, low-metallicity ($Z=0.001$), Salpeter-IMF galaxy experiencing a short burst of star formation ($10^8$ yr) and passive evolution thereafter (black line). The black triangle indicates an age of $10^6$ yr and the black square $5\\times 10^8$ yr (roughly the age of the Universe at this redshift). The star symbols indicate the colours of two cool, $z=10$ dark stars from Table~\\ref{modeltable}: the 7500 K, $690 \\ M_\\odot$ model from the 1 GeV WIMP track (blue star) and the 5800 K, $106 M_\\odot$ model from the 100 GeV WIMP track (red star). Both of these (as is the case for all $T_\\mathrm{eff}<10000$ K dark stars observed in these filters; see Fig.~\\ref{colour_criteria}) are considerably redder than the colours expected for galaxies, regardless of their age. Dashed lines indicate how the colours of the model galaxy (at an assumed age of $10^8$ yr) would shift if it were to contain dark stars of either of the two types. The filled circles along the dashed tracks indicate mixtures at which these dark stars make up $1\\%$ of the stellar mass in the model galaxy. These points are also significantly redder than the reddest point along the galaxy track, indicating that a $\\sim 1\\%$ stellar mass fraction in dark stars within $z\\approx 10$ galaxies would be detectable through multiband photometry. \\label{colour_criteria2}} \\end{figure}" }, "1002/1002.3619_arXiv.txt": { "abstract": "We complete the census of nuclear X-ray activity in 100 early type Virgo galaxies observed by the \\cxo\\ {\\it X-ray Telescope} as part of the AMUSE-Virgo survey, down to a (3$\\sigma$) limiting luminosity of $3.7\\times 10^{38}$ \\es\\ over 0.5-7 keV. The stellar mass distribution of the targeted sample, which is mostly composed of formally `inactive' galaxies, peaks below $10^{10}$ \\msun, a regime where the very existence of nuclear super-massive black holes (SMBHs) is debated. Out of 100 objects, 32 show a nuclear X-ray source, including 6 hybrid nuclei which also host a massive nuclear cluster as visible from archival \\hst\\ {\\it Space Telescope} images. After carefully accounting for contamination from nuclear low-mass X-ray binaries based on the shape and normalization of their X-ray luminosity function, we conclude that between $24-34\\%$ of the galaxies in our sample host a X-ray active SMBH (at the 95$\\%$ C.L.). This sets a firm lower limit to the black hole occupation fraction in nearby bulges within a cluster environment. The differential logarithmic X-ray luminosity function of active SMBHs scales with the X-ray luminosity as \\lx$^{-0.4\\pm0.1}$ up to $10^{42}$ erg \\se. At face value, the active fraction --down to our luminosity limit-- is found to increase with host stellar mass. However, taking into account selection effects, we find that the average Eddington-scaled X-ray luminosity scales with black hole mass as \\bhm$^{-0.62^{+0.13}_{-0.12}}$, with an intrinsic scatter of $0.46^{+0.08}_{-0.06}$ dex. This finding can be interpreted as observational evidence for `down-sizing' of black hole accretion in local early types, that is, low mass black holes shine relatively closer to their Eddington limit than higher mass objects. As a consequence, the fraction of active galaxies, defined as those above a fixed X-ray Eddington ratio, {\\it decreases} with increasing black hole mass. ", "introduction": "\\label{sec:intro} Historically speaking, active galaxies are characterized by compact nuclei with abnormally high luminosity and fast variability ascribed to accretion of mass onto a super-massive black hole (SMBH). While the term AGN (active galactic nuclei) generally refers to nuclear luminosities in excess of $10^{43-44}$ \\es, the distinction between active and inactive is rather arbitrary, that is, is set by our ability to detect and interpret signatures of accretion-powered activity. From elaboration of the Soltan argument (Soltan 1982) follows that, since black holes have grown mostly via radiatively efficient accretion as powerful quasars (e.g. Yu \\& Tremaine 2002; Marconi \\etal 2004; Merloni \\& Heinz 2008; Shankar, Weinberg \\& Miralda-Escud\\'e\\ 2009), nearby galaxies ought to harbor, if anything, only weakly accreting black holes. The alleged ubiquity of SMBHs at the center of (massive) galaxies, together with the realization that BHs play a crucial role in regulating the assembly history and evolution of their hosts (Kormendy \\& Richstone 1995; Kormendy \\etal 1997; Magorrian et al. 1998; Gebhardt et al. 2000; Ferrarese \\& Merritt 2000; McLure \\& Dunlop 2002; Marconi \\& Hunt 2003; Ferrarese \\& Ford 2005; G\\\"ultekin \\etal 2009), have spurred a series of searches for active nuclei in the nearby universe at different wavelengths, each with its own advantages and limitations. In particular, the low mass end of the black hole mass function in the local universe (Greene \\& Ho 2007, 2009) remains poorly constrained, and can only be explored indirectly, as even the highest angular resolution attainable with current instrumentation is not sufficient to go after a $\\simlt 10^6$ \\msun\\ black hole through resolved stellar kinematics, except for exceptionally nearby systems. (see Bentz \\etal 2009 on nearby, reverberation-mapped AGN). Amongst the general population of galaxies, optical studies suggest that nuclear activity is quite common (43$\\%$ of all galaxies in the Palomar sample; Ho, Filippenko \\& Sargent 1997; Ho 2008). The percentage raises substantially in galaxies with a prominent bulge component, approaching 70$\\%$ for Hubble types E-Sb. The dependence of the nuclear properties on Hubble type -- with late type objects displaying active fractions as low as $10\\%$ -- has been confirmed by numerous other studies (e.g. Kauffmann et al. 2003; Miller et al. 2003; Decarli et al. 2007), although a recent work based on high-resolution mid-infrared spectrometry of a sample of (32) inactive galaxies, suggests that the AGN detection rate in late-type galaxies is possibly 4 times larger than what optical observations alone indicate (Satyapal \\etal 2008). Since the enormous wealth of data from the Sloan Digital Sky Survey (SDSS) became available, various environmental effects on nuclear activity have been investigated, such as host galaxy properties (Kauffmann \\etal 2003, 2004; Rich et al. 2005; Kewley \\etal 2006; Schawinski et al. 2007, 2009; Kauffmann, Heckman \\& Best 2008) local density and large scale environment (Kauffmann et al. 2003, 2004; Constantin \\& Vogeley 2006; Constantin et al. 2008; Choi \\etal 2009). While AGN were first discovered as powerful, unresolved optical sources at the center of galaxies, emission at higher frequencies, hard X-rays and gamma-rays, is almost univocally associated with non thermal processes related to accretion, such as Comptonization of thermal photons in a hot electron-positron plasma. Hard X-rays in particular offer a clean-cut diagnostics, and a relatively unexplored one, to pinpoint low accretion power SMBHs in nearby galaxies. So far, searches for nuclear X-ray sources in formally inactive galaxies have been somewhat sparse and focused on the high-mass end of the local population. Prior to the launch of \\cxo\\ and {\\it XMM-Newton}, such observations were effectively limited to X-ray luminosities $\\simgt 10^{40}$ erg \\se\\ even in the nearest elliptical galaxies (e.g.: Canizares, Fabbiano \\& Trinchieri 1987; Fabbiano \\etal 1993; Fabbiano \\& Juda 1997; Allen, Di Matteo \\& Fabian 2000). Due to the lack of sensitivity and angular resolution, earlier mission were necessarily deemed to confusion between accretion-powered sources of various nature, most notably nuclear vs. off-nuclear, and also thermally emitting gas. As an example, Roberts \\& Warwick (2000), report on the detection of 54 X-ray cores out of 83 Palomar galaxies targeted by {\\it ROSAT}. As also noted by Ho (2008), a significant (dominant, we argue) fraction of the cores' X-ray flux may be due to unresolved emission from X-ray binaries. The greatly improved sensitivity and angular resolution of \\cxo~and {\\it XMM-Newton} have made it possible to investigate nuclear emission associated with SMBHs orders of magnitude deeper, effectively bridging the gap between AGN and inactive galaxies (e.g. Di Matteo et al. 2000, 2001, 2003; Ho et al. 2001; Loewenstein et al. 2001; Sarazin, Irwin \\& Bregman 2001; Fabbiano \\etal 2003, 2004; Terashima \\& Wilson 2003; Pellegrini 2005; Soria et al. 2006a, 2006b; Santra et al. 2007; Pellegrini, Ciotti \\& Ostriker 2007; Ghosh \\etal 2008; Zhang \\etal 2009). Direct measurements of bolometric Eddington ratios in {\\it bona fide} AGN are typically no lower than $10^{-3}$ (e.g. Woo \\& Urry 2002; Kollmeier \\etal 2006; Heckman \\etal 2004). As a comparison, the inferred Eddington-scaled X-ray luminosities of inactive galaxies -- that is, of their nuclear SMBHs -- are as low as $10^{-8}$: in those massive elliptical galaxies where the temperature and density profiles of the thermally-emitting gas can be reconstructed and used to estimate the inner gas reservoir available to accretion, the measured nuclear X-ray luminosities are orders of magnitude lower than expected from Bondi-type accretion onto the nuclear SMBH (e.g. Pellegrini 2005; Soria \\etal 2006a,b). However, most X-ray studies target massive nearby elliptical galaxies, and are thus biased towards the high mass end of the SMBH mass function. In order to expand our knowledge about black hole demographics in the local universe, it is necessary to explore both the low mass {and} the low-luminosity end of the distribution. Even in the nearby universe, pushing the threshold down to X-ray luminosities as low as a few $10^{38}$ erg \\se\\ necessarily means facing contamination from bright X-ray binaries within the instrument point spread function (PSF). This problem has been touched upon in a recent work by Zhang \\etal (2009), who collected archival \\cxo\\ observations of 187 galaxies (both late and early types) within 15 Mpc. 86 of them host nuclear X-ray cores, the majority of which, based on the fitted slope of their differential luminosity function, are attributed to low-level accretion onto SMBHs, rather than to X-ray binaries. The issue of X-ray binary contamination becomes particularly delicate when the inferred X-ray active fractions are then used to place constraints on the local black hole occupation fraction. From an observational standpoint, the very existence of SMBHs in nearby dwarf galaxies remains a matter of investigation. Ferrarese \\etal (2006a) argue that the creation of a `central massive object', be it a black hole or a compact stellar nucleus, would be the natural byproduct of galaxy evolution, with the former being more common in massive bright galaxies (with absolute B magnitude M$_{\\rm B}$ brighter than $-$20), and the latter dominating --possibly taking over-- at magnitudes fainter than $-18$ (see also Wehner \\& Harris 2006, and Kormendy \\etal 2009). Massive nuclear star clusters (e.g. Seth \\etal 2008, 2010; Graham \\& Spitler 2009), with inferred radii around a few tens of pc, become increasingly prominent down the mass function. When dealing with faint X-ray cores ($10^{38}-10^{39}$ erg \\se), the problem of X-ray binary contamination is further exacerbated by the presence of a nuclear star cluster, having higher stellar encounter rates, and hence a higher X-ray binary fraction with respect to the field. In order to deliver an unbiased census of nuclear activity for nearby galaxies down to X-ray luminosities as low as the Eddington limit for a solar mass object, not only it is mandatory to deal with nearby sources, but it also becomes necessary to have information about their stellar content within the X-ray instrument PSF, specifically the presence/absence of a nuclear star cluster. Additionally, in order to avoid contamination from the short lived, X-ray bright, high-mass X-ray binaries, deep X-ray searches for weakly active SMBHs (down to $\\sim 10^{38}$ erg \\se) should be preferentially limited to the nuclei of early type galaxies. \\\\ To this purpose, and with these caveats in mind, in Cycle 8 we proposed and were awarded a large \\cxo /\\spi\\ program to observe 100 (84 new+16 archival) spheroidal galaxies in the Virgo Cluster (AMUSE-Virgo: PI: Treu, 454 ks). The targeted sample is that of the \\hst\\ Virgo Cluster Survey (VCS; C\\^ot\\'e \\etal 2004). For each galaxy, the high resolution $g$ and $z$ band images enable us to resolve, when present, the nuclear star cluster, infer its enclosed mass (following Ferrarese \\etal 2006a; see Ferrarese \\etal 2006b, for a detailed isophotal analysis of the \\hst\\ data), and thus estimate the chance contamination from a low-mass X-ray binary (LMXB) as bright/brighter than the detected X-ray core based on the shape and normalization of the LMXB luminosity function in external galaxies (see \\S\\ref{sec:lmxb}, and references therein). As a part of the AMUSE-Virgo survey, each VCS galaxy was observed with \\cxo\\ for a minimum of 5.4 ks, which, at the average distance of Virgo (16.5 Mpc; Mei \\etal 2007; see also Blakeslee et al. 2009), yields a (3$\\sigma$) sensitivity threshold of $3.75\\times 10^{38}$ erg \\se\\ over the \\cxo\\ bandpass. The \\cxo\\ results from the first 32 targets (16 new + 16, typical more massive, archival observations) have been presented by Gallo \\etal (2008; hereafter Paper I.). Point-like X-ray emission from a position coincident with the optical nucleus was detected in 50 per cent of the galaxies. We argued that, for this sub-sample, all of the detected nuclear X-ray sources are most likely powered by low-level accretion on to a SMBH, with a $\\simlt$11 per cent chance contamination from LMXBs in one case only (VCC1178=NGC 4486B, for which independent evidence points towards the presence of a nuclear BH; Lauer \\etal 1996). The incidence of nuclear X-ray activity increases with the stellar mass \\mstar\\ of the host galaxy: only between 3--44$\\%$ of the galaxies with \\mstar$<10^{10}$ \\msun\\ harbor an X-ray active SMBH. The fraction rises to between 49--87$\\%$ in galaxies with stellar mass above $10^{10}$ \\msun. (at the 95$\\%$ C.L.). In Paper II. we complete the X-ray analysis of the whole AMUSE-Virgo sample: the final deliverable product of this study is an unbiased census of accretion-powered luminosity in a galaxy cluster environment, and thus the first measurement of the SMBH activity duty cycle. The Paper is organized as follows: \\S~\\ref{sec:data} summarizes our analysis of the new \\cxo\\ data and the stacking procedure. In \\S~\\ref{sec:lmxb} we carefully address the issue of contamination from low-mass X-ray binaries to the detected X-ray cores. \\S~\\ref{sec:res} and 5 presents our main results, specifically on the active fraction, dependence of accretion luminosity on black hole mass and X-ray luminosity function. We discuss the implications of our results in \\S~\\ref{sec:disc}, and conclude with summary in~\\S~\\ref{sec:conc}. We refer the reader to Paper I. for a thorough description of the program, as well as the determination of the parameters (such as stellar masses, stellar velocity dispersions, and black hole masses) employed throughout this series. In a companion paper, we report on the results from the \\spi\\ 24 $\\mu$m observations of the same sample (Leipski \\etal, Paper III.). ", "conclusions": "\\label{sec:disc} \\subsection{Super-massive black holes vs. massive nuclear star clusters} Until recently it was generally believed that massive black holes and nuclear star clusters did not generally coexist at the centres of galaxies. Less than a handful of counter-examples (e.g. Filippenko \\& Ho 2003; Graham \\& Driver 2007) were the exceptions to confirm the rule. More systematic studies, e.g. by Ferrarese \\etal (2006a) and Wehner \\& Harris (2006), showed the transition between galaxies which host predominately a black hole vs. a nuclear cluster occurs around 10$^{10}$ \\msun. However, while the latter conclude that nucleated galaxies show no evidence of hosting SMBHs, the former speculate that nuclei form in all galaxies but they are destroyed by the evolution of preexisting SMBHs or collapse into a SMBH in the most massive cases. It is also suggested that ``SMBHs and nuclei are almost certainly mutually exclusive in the faintest galaxies belonging to the VCS sample\\footnote{The same considered in this work.}\". This is indicated by the fact that, although the nuclear masses of NGC~205 and M33 are fully consistent with the relation\\footnote{See Figure 2 in Ferrarese \\etal (2006a).} between the mass the nuclear object (be it a cluster or a BH) and the host galaxy, the upper limits on their SMBH masses are not, implying that {`` [..] neither galaxy contains an SMBH of the sort expected from extrapolations of the scaling relations defined by SMBHs in massive galaxies\"}. Merritt (2009) examines the evolution of nuclear star clusters with and without SMBHs from a theoretical point of view, finding that nuclear star clusters with black holes are always bound to expand, due primarily to heating from the galaxy and secondarily to heating from stellar disruptions. As a consequence, core-collapsed clusters should not be harboring nuclear black holes. From an observational point of view, there has been a number of efforts to quantify the degree of coexistence of nuclear clusters and SMBHs, and their mutual properties. Seth et al.\\ (2008) searched for active nuclei in 176 galaxies with known nuclear clusters, using optical spectroscopy, X-ray and radio data. They find that the AGN fraction increases strongly with increasing galaxy and nuclear cluster mass, consistent with previous studies of the general galaxy population. In addition, the variation of the AGN fraction with Hubble type is also consistent with the whole Palomar sample (Ho \\etal 1997), indicating that the presence (or absence) of a nuclear star cluster does not play a crucial role in boosting (or hampering) accretion-powered activity onto a SMBH (see also Gonzalez-Delgado et al.\\ 2008). More recently, Graham \\& Spitler (2009) reported on 12 new systems which host both a nuclear star cluster and a SMBHs, and for which they were able acquire both the masses of the nuclear components, as well as the stellar mass of the host spheroid. They find that, for host stellar masses in the range $10^{8-11}$ \\msun, the nucleus-to-spheroid mass ratio decreases from a few to about 0.3$\\%$. This ratio is expected to saturate to a constant value once dry merging commences, and the nuclear cluster disappear. Our work tackles the issue of nuclear SMBH-star cluster coexistence taking a complementary approach with respect to that of Seth et al., that is, we ask how many of those nuclei which show a nuclear cluster also host a (X-ray active) SMBH. We remind the reader that the VCS sample (C\\^ot\\'e \\etal 2004) surveyed by AMUSE-Virgo is complete down to a $B$ magnitude of $-$18, and is a random sample of fainter (early type) objects. In terms of (hosts') stellar mass distribution, the sample peaks well below $10^{11}$ \\msun, and is thus particularly well suited for investigating hybrid potentially nuclei. While 32 out 100 galaxies are found to host a nuclear X-ray source, only 6 of them also show evidence for a nuclear star cluster as visible from archival \\hst~images (typically, the star clusters are identified as overdr-densities above a single-component Sersic profile, but see Ferrarese \\etal 2006b for details about the fitting procedure). After taking into account LMXB contamination, while the fraction of hybrid nuclei as function of host stellar mass \\mstar\\ is constrained between 0.3 and 7$\\%$ for log\\mstar$>11$ (95$\\%$ C.L., and down to a limiting 2-10 keV luminosity of $\\sim 2\\times 10^{38}$ \\es), the lack of star cluster--SMBH matches above $10^{11}$ sets an upper limit of 32$\\%$ to the fraction of such hybrid nuclei in massive early types. \\subsection{Active fraction and down-sizing in black hole accretion} The bottom panel of Figure~\\ref{fig:activefrac} shows how he fraction $f_{X}$ of objects hosting an active SMBH --down to our luminosity limit-- increases as a function of the host mass \\mstar: $0.01{10.5}$, and $0.166$ have been discovered using {\\em Swift}: GRB 050904 at $z=6.3$ in 2005 \\citep{Cusumano_etal:2006, Kawai_etal:2006}; GRB 080913 at $z=6.7$ in 2008 \\citep{Greiner_etal:2009}; and, in 2009, the remarkable source GRB 090423 at $z=8.2$ \\citep{Tanvir_etal:2009,Salvaterra_etal:2009}, the most distant object yet found. \\cite{Totani_etal:2006} present a detailed analysis of the optical afterglow spectrum of GRB 050904, which was measured with the highest S/N of the three $z>6$ GRBs. The spectrum displays a red damping wing from Ly$\\alpha$ absorption that is produced by some combination of a high-column density absorber near the GRB (hereafter referred to as a DLA, for damped Ly$\\alpha$ absorber), and a smoothly distributed component in the IGM. By fitting absorption models containing these two components, \\cite{Totani_etal:2006} were able to place a limit on the neutral fraction of the IGM at $z=6.3$ of $\\xhi<0.17(0.60)$ at $68\\%(95\\%)$ confidence. The constraints are relatively weak despite the reasonably high S/N of the spectrum. The sources GRB 080913 and GRB 090423 are potentially more interesting, because of their higher redshifts, but the published spectra have insufficient S/N to place any useful constraints on $\\xhi$ when employing two-component (DLA+IGM) fits. Here we describe and analyse an improved spectrum of GRB 080913, which includes unpublished spectroscopic data taken three nights after the published spectrum. In Section 2 we describe the observations taken on each night, and the reduction techniques. We analyse the combined spectrum in Section 3, first searching for absorption lines in the new spectrum, in order to measure the redshift of the source, and then fitting a two-component DLA+IGM model to the observed continuum break. Finally, the results are summarised in Section 4. ", "conclusions": "New optical spectroscopic observations of GRB 080913 have been presented and analysed. The detection of SII+SiII absorption ($0.1260\\mu$m) at $2.9\\sigma$ provides a redshift of the DLA host galaxy of $z=6.733$. Employing a joint DLA$+$IGM model to fit the observed continuum break, we find an upper limit to the neutral fraction of the IGM $\\xhi<0.73$ at a probability of $90\\%$. However this result rests on the assumption that the ionised region surrounding the host galaxy is transparent to Ly$\\alpha$. Any analysis of a GRB spectrum needs to include the radius of the ionised region as a free parameter, and to consider the question of the neutral fraction within this zone. Furthermore the scatter in measurements of $\\xhi$ between different sources at similar redshifts is predicted to be substantial \\citep{McQuinn_etal:2008,Mesinger_Furlanetto:2008}, and needs to be quantified. Higher S/N spectra of several sources at high redshift will be required to make significant progress in this field." }, "1002/1002.3872_arXiv.txt": { "abstract": "For a wide variety of initial and boundary conditions, adiabatic one dimensional flows of an ideal gas approach self-similar behavior when the characteristic length scale over which the flow takes place, $R$, diverges or tends to zero. It is commonly assumed that self-similarity is approached since in the $R\\rightarrow\\infty(0)$ limit the flow becomes independent of any characteristic length or time scales. In this case, the flow fields $f(r,t)$ must be of the form $f(r,t)=t^{\\alpha_f}F(r/R)$ with $R\\propto(\\pm t)^\\alpha$. We show that requiring the asymptotic flow to be independent only of characteristic length scales implies a more general form of self-similar solutions, $f(r,t)=R^{\\delta_f}F(r/R)$ with $\\dot{R}\\propto R^\\delta$, which includes the exponential ($\\delta=1$) solutions, $R\\propto e^{t/\\tau}$. We demonstrate that the latter, less restrictive, requirement is the physically relevant one by showing that the asymptotic behavior of accelerating blast-waves, driven by the release of energy at the center of a cold gas sphere of initial density $\\rho\\propto r^{-\\omega}$, changes its character at large $\\omega$: The flow is described by $0\\le\\delta<1$, $R\\propto t^{1/(1-\\delta)}$, solutions for $\\omega<\\omega_c$, by $\\delta>1$ solutions with $R\\propto (-t)^{1/(\\delta-1)}$ diverging at finite time ($t=0$) for $\\omega>\\omega_c$, and by exponential solutions for $\\omega=\\omega_c$ ($\\omega_c$ depends on the adiabatic index of the gas, $\\omega_c\\sim8$ for $4/3<\\gamma<5/3$). The properties of the new solutions obtained here for $\\omega\\ge\\omega_c$ are analyzed, and self-similar solutions describing the $t>0$ behavior for $\\omega>\\omega_c$ are also derived. ", "introduction": "\\label{sec:introduction} Self-similar solutions to the hydrodynamic equations describing adiabatic one dimensional flows of an ideal gas are of interest for several reasons. The non-linear partial differential hydrodynamic equations are reduced for self-similar flows to ordinary differential equations, which greatly simplifies the mathematical problem of solving the equations and in certain cases allows one to find analytic solutions. Moreover, self-similar solutions often describe the limiting behavior approached asymptotically by flows which take place over a characteristic scale, $R$, which diverges or tends to zero \\citep[see][for reviews]{SedovBook,ZelDovichRaizer,BarenblattBook}. Some examples of such asymptotic solutions which are widely used in astrophysical contexts are the Sedov-von Neumann-Taylor solutions \\citep{Sedov46,vonNeumann47,Taylor50} describing expanding decelerating spherical blast waves, for which $R\\rightarrow\\infty$, and the Gandel'Man-Frank-Kamenetskii--Sakurai solutions \\citep{GandelMan56,Sakurai60} describing the emergence of a shock wave from the surface of a star, for which $R\\rightarrow0$. Both types of solutions are relevant, e.g., to supernova explosions and in particular to the recently detected shock breakouts \\citep[e.g.][]{MM99,WaxmanMeszarosCampana07}. An extensive discussion of spherical self-similar blast-waves in an astrophysical context is given by \\citet{OstrikerMcKee88}. It is commonly assumed that self-similarity is approached since in the $R\\rightarrow\\infty(0)$ limit the flow becomes independent of any characteristic length or time scales. Using dimensional arguments it is possible to show that if the flow is determined by a set of constants, using which it is impossible to construct a constant with the dimensions of length or time, then the flow fields $f(r,t)$ must be of the form $f(r,t)=t^{\\alpha_f}F(r/R)$ with $R\\propto(\\pm t)^\\alpha$ \\citep[see chapter XII of][]{ZelDovichRaizer}. We show in \\S~\\ref{sec:general} that requiring the asymptotic flow to be independent only of characteristic length scales is sufficient for showing, based on dimensional arguments, that the flow must be self-similar. The less restrictive requirement allows a more general form of self-similar solutions, $f(r,t)=R^{\\delta_f}F(r/R)$ with $\\dot{R}\\propto R^\\delta$, which includes the exponential ($\\delta=1$) solutions, $R\\propto e^{t/\\tau}$. The existence of exponential self-similar solutions has been noted by several authors \\citep[e.g.][]{StanyukovichBook}. However, it was generally assumed that asymptotic solutions, which are of interest, are of a power law form. In \\S~\\ref{sec:blast} we show that the asymptotic self-similar solutions describing the propagation of accelerating blast waves, propagating in a cold gas sphere of initial density $\\rho\\propto r^{-\\omega}$ with $\\omega>3$, are of the more general form, $\\dot{R}\\propto R^\\delta$, with exponential solutions obtained at $\\omega=\\omega_c(\\gamma)$, where $\\gamma$ is the adiabatic index of the gas. The new solutions obtained here for $\\omega\\ge\\omega_c$ extend the family of second type solutions describing the asymptotic flow of accelerating blast waves, which was derived by \\citet{WaxmanShvarts93} and was limited to $\\omega<\\omega_c$, to $\\omega\\ge\\omega_c$. The properties of the new solutions are analyzed in \\S~\\ref{sec:WS}, and self similar solutions describing the $\\omega>\\omega_c$ flow at times later than the finite divergence time are derived in \\S~\\ref{sec:t0}. Our results are summarized in \\S~\\ref{sec:summary}. It should be noted here that the self-similar solutions derived by \\citet{WaxmanShvarts93} exist only for $\\omega>\\omega_g(\\gamma)>3$, where $\\omega_g=3.26$ for $\\gamma=5/3$ and approaches 3 for $\\gamma\\rightarrow1$, while the Sedov-von Neumann-Taylor solutions provide the correct asymptotic solutions only for $\\omega<3$. The asymptotic behavior within the (narrow) range of $3<\\omega<\\omega_g(\\gamma)$ is not described by either of the two types of solutions. The nature of the asymptotic flow in this regime is discussed in \\citet{Gruzinov03,KushnirGap}. ", "conclusions": "\\label{sec:summary} We have discussed the asymptotic behavior of adiabatic one dimensional flows of an ideal gas, which take place over a characteristic scale $R$ that diverges or tends to zero. We have shown in \\S~\\ref{sec:general} that requiring the asymptotic flow to be independent of characteristic length scales implies, based on dimensional arguments, that the flow must be self-similar, with flow fields given by eq.~(\\ref{eq:ss_scaling}) and $R$ satisfying eq.~(\\ref{eq:Rdot}), $\\dot{R}\\propto R^\\delta$. The ordinary differential equations determining the self-similar profiles are given by eqs.~(\\ref{eq:dUdC})--(\\ref{eq:deltas}). In \\S~\\ref{sec:blast} we have shown that the asymptotic self-similar solutions describing the propagation of accelerating blast waves, propagating in a cold gas sphere of initial density $\\rho\\propto r^{-\\omega}$ with $\\omega>3$, are of the general form $\\dot{R}\\propto R^\\delta$, with exponential solutions obtained at $\\omega=\\omega_c(\\gamma)$. $\\delta(\\omega,\\gamma)$ is shown in fig.~(\\ref{fig:delta}) for $\\gamma=4/3,5/3$. Fig.~\\ref{fig:UC} demonstrates that numerical solutions of the hydrodynamic equations, eq.~(\\ref{eq:hydro_eq}), indeed approach these self-similar solutions as $R$ diverges. The properties of the self-similar solutions obtained for $\\omega\\ge\\omega_c(\\gamma)$ are analyzed in \\S~\\ref{sec:WS}. The $C(U)$ curves of these solutions approach the singular point $\\{U=0,C=0\\}$ as $\\xi\\rightarrow0$. Analyzing the behavior of the solutions near this singular point we have shown that the energy and mass contained in region of the $\\{\\xi,R\\}$ plane bounded by $\\xi=1$ and $\\xi=\\xi_+(R)$, where $\\xi_+(R)$ is a $C_+$ characteristic that approaches $\\xi=0$ as $R$ diverges, both tend to finite constants as $R$ diverges. For $\\delta>1$, the shock radius diverges at a finite time, $t=0$. This implies that the spatial distribution of the flow fields at the divergence time is described by the $\\xi\\rightarrow0$ behavior of the solution, which is given by eq.~(\\ref{eq:asym_dg1}). This implies that at $t=0$ the spatial distribution of the flow fields is given by eq.~(\\ref{eq:t0}). In \\S~\\ref{sec:t0} we have shown that the $t>0$ flow is described in this case by a self-similar solution with the same value of $\\delta$ as that of the $t<0$ solution, see eq.~(\\ref{eq:R_tilde}), and spatial profiles determine by eq.~(\\ref{eq:t0_prof})." }, "1002/1002.5045_arXiv.txt": { "abstract": "I show how prior work with R. Wald on geodesic motion in general relativity can be generalized to classical field theories of a metric and other tensor fields on four-dimensional spacetime that 1) are second-order and 2) follow from a diffeomorphism-covariant Lagrangian. The approach is to consider a one-parameter-family of solutions to the field equations satisfying certain assumptions designed to reflect the existence of a body whose size, mass, and various charges are simultaneously scaled to zero. (That such solutions exist places a further restriction on the class of theories to which our results apply.) Assumptions are made only on the spacetime region outside of the body, so that the results apply independent of the body's composition (and, e.g., black holes are allowed). The worldline ``left behind'' by the shrinking, disappearing body is interpreted as its lowest-order motion. An equation for this worldline follows from the ``Bianchi identity'' for the theory, without use of any properties of the field equations beyond their being second-order. The form of the force law for a theory therefore depends only on the ranks of its various tensor fields; the detailed properties of the field equations are relevant only for determining the charges for a particular body (which are the ``monopoles'' of its exterior fields in a suitable limiting sense). I explicitly derive the force law (and mass-evolution law) in the case of scalar and vector fields, and give the recipe in the higher-rank case. Note that the vector force law is quite complicated, simplifying to the Lorentz force law only in the presence of the Maxwell gauge symmetry. Example applications of the results are the motion of ``chameleon'' bodies beyond the Newtonian limit, and the motion of bodies in (classical) non-Abelian gauge theory. I also make some comments on the role that scaling plays in the appearance of universality in the motion of bodies. ", "introduction": "In special relativity, a non-interacting body moves in a straight line. Therefore, it is not surprising that in general relativity an ``infinitesimal test body'' (i.e., a body small enough that the curvature of the external universe can be neglected, and weakly-gravitating enough that curvature it generates can be neglected) will move \\textit{locally} in a straight line, i.e., it will follow a geodesic. But from this perspective it does seem quite surprising that strong-field bodies like neutron stars and black holes in fact \\textit{also} move on geodesics (in the limit of small size). After all, no matter how small or light such a body, the local spacetime metric will differ significantly from that of flat spacetime, and one would therefore expect that nonlinear gravitational dynamics---certainly not special relativity---would principally determine its motion. Furthermore, since the metrics of different strong-field bodies will differ greatly from \\textit{each other}, one would perhaps expect there to be \\textit{no} universal law for the motion of strong-field bodies at all. Indeed, the natural assumption would seem to be that the motion of a strong-field body depends in detail upon its composition. This expectation is incorrect for a very counter-intuitive reason: in general relativity, the motion of a small body is in fact completely determined by field dynamics \\textit{outside} of the body. This surprising fact was first demonstrated by Einstein, Infeld and Hoffman \\cite{einstein-infeld-hoffman}, and has become the foundation of a more modern approach to motion termed ``matched asymptotic expansions'' \\cite{matched-expansions} (see also \\cite{poisson,gralla-wald,futamase-hogan-itoh,pound}). The basic physical requirement of this line of work is the existence of a region (the ``buffer zone'') sufficiently far from the body that the body field may be approximated as a multipole series, yet sufficiently close to the body that the field of the external universe may be approximated in an ordinary Taylor series. The vacuum gravitational dynamics taking place in this region then suffice to determine the motion. A primary purpose of this paper is to determine to what extent this conclusion generalizes to other classical field theories. To investigate this question I generalize the approach taken in \\cite{gralla-wald} to deriving geodesic motion in general relativity.\\footnote{I do not treat self-force corrections, which were the primary focus of \\cite{gralla-wald}.} In the formalism of \\cite{gralla-wald} a small body is characterized by a one-parameter-family of solutions to the vacuum Einstein equation describing the region outside of a body that shrinks to zero size and mass with the perturbation parameter, $\\lambda$. A family with such behavior is considered by demanding the existence of a second, ``scaled'' limit wherein the coordinates and metric are rescaled such the body is held at fixed size and mass. At $\\lambda=0$ in the original limit the body disappears, leaving behind a smooth spacetime with a preferred worldline, $\\gamma$, picked out; this worldline is interpreted as the lowest-order perturbative motion of the body. We showed that $\\gamma$ must be a geodesic by applying the Bianchi identity to an effective point particle description that (remarkably) emerges at first order in $\\lambda$. In this paper I generalize the approach to theories that 1) follow from a diffeomorphism-covariant Lagrangian, ensuring a ``Bianchi identity'' and 2) have second-order field equations. For the above class of theories the method of \\cite{gralla-wald} gives an equation for $\\gamma$ that depends only on ``buffer zone'' field properties, showing that the Einstein-Infeld-Hoffman idea remains correct in a more general context. More specifically, the equation involves, in addition to the value and first-derivative of the external fields at the location of the body, various charges (understood to include mass as the charge associated with the metric) that are determined from the body's fields in the scaled limit (they are ``field monopoles''). The results rely only on properties 1) and 2) above and are therefore surprisingly independent of the details of the theory. In particular, the force law depends only on the form of the Bianchi identity, which in turn depends only on the ranks of the tensor fields considered (although extra identities following from gauge symmetry can greatly simplify the results). Therefore, the expression for the force in terms of the charges and external field values is in fact identical across theories with the same types of tensor fields and the same gauge symmetries. However, the charges associated with a particular body composition will differ in different theories, since the relationship between a given source and the field monopoles it generates will depend on the field equations. In this way different theories will make different predictions for the motion of the ``same body,'' even when the force law is identical. In interpreting the results it is useful to distinguish varying degrees of ``universality'' in the motion of small bodies. In the case of general relativity, all small bodies move on geodesics, so that their internal structure is completely irrelevant to their motion. In Einstein-Maxwell theory, a single number characterizing the body (the charge-to-mass ratio) determines how it will move, so that the internal structure is minimally relevant. In scalar-tensor theory, a free function of time (the charge-to-mass ratio of the non-conserved charge) specifies the motion of a body, so that the internal structure is somewhat relevant. In higher-rank theories a finite number of free functions of time characterize the motion of a body. Of these results only geodesic motion in general relativity is truly universal in that it applies to \\textit{all} bodies; however, I will refer to all of the above results as ``universal behavior in motion'', since the information required to determine the motion of a small body is reduced from the complete description of the body to the knowledge of a finite number of parameters at each time. To adopt the language of condensed matter physics, there are thus large ``universality classes'' of small bodies that move in the same way. The content of this paper is as follows. In section \\ref{sec:GR} I summarize the formalism of \\cite{gralla-wald} to derive geodesic motion in general relativity. In section \\ref{sec:scalar} I generalize the formalism to Einstein-scalar and then more general scalar-tensor theories, deriving the scalar force law. Note that mass evolution always occurs, and the scalar charge evolution is unconstrained. I discuss the results in the context of specific scalar-tensor theories and comment on scaling and universality. In section \\ref{sec:vector} I apply the formalism to vector-tensor theories to derive the vector force law. This surprisingly complicated equation simplifies to the Lorentz force law in theories with the Maxwell gauge symmetry. I also derive the simplified force law in the case of non-Abelian gauge theory. Finally in section \\ref{sec:general} I give the proof that universality in motion is achieved via buffer zone dynamics for tensor fields of arbitrary rank. A definition and disambiguation of scale-invariance is given in an appendix. I use the conventions of Wald \\cite{wald} and work in units where $G=c=1$. Early-alphabet Latin indices $a,b,...$ are abstract spacetime indices, while Greek indices $\\mu,\\nu,...$ give tensor components in a coordinate system. When working in coordinates $(t,x^1,x^2,x^3)$, mid-alphabet Latin indices $i,j,...$ denote spatial components $1-3$, while a zero denotes the time component $t$. Mid-alphabet capital Latin indices $I,J,...$ label members of a collection of tensor fields. ", "conclusions": "I have treated the motion of small bodies in classical field theory via the approach of \\cite{gralla-wald}. The search for one-parameter-families of solutions representing the exterior field of a shrinking body led precisely to the physical assumption of a ``buffer zone''---a region far enough from the body that its field can be approximated in a multipole series, but close enough to the body that the field of the external universe can be approximated in an ordinary Taylor series. No assumptions about the body interior are made. In the case of second-order metric-based theories following from a diffeomorphism-covariant Lagrangian, I derived the force law for scalar and vector fields, and showed that the method works in the general-rank case. This provides a rigorous derivation of the small-body force law in many classical field theories commonly considered, and shows that field dynamics outside a body determines its motion in a very general class of theories." }, "1002/1002.0456_arXiv.txt": { "abstract": "Using the MegaCam imager on the Canada-France-Hawaii Telescope, we have resolved individual stars in the outskirts of the nearby large spiral galaxy M81 (NGC~3031) well below the tip of the red giant branch of metal-poor stellar populations over $\\sim 60 \\kpc \\times 58 \\kpc$. In this paper, we report the discovery of new young stellar systems in the outskirts of M81. The most prominent feature is a chain of clumps of young stars distributed along the extended southern H{\\sc i} tidal arm connecting M~81 and NGC~3077. The colour-magnitude diagrams of these stellar systems show plumes of bright main sequence stars and red supergiant stars, indicating extended events of star formation. The main sequence turn-offs of the youngest stars in the systems are consistent with ages of $\\sim 40$ Myr. The newly reported stellar systems show strong similarities with other known young stellar systems in the debris field around M81, with their properties best explained by these systems being of tidal origin. ", "introduction": "\\label{intro} \\footnotetext[1]{This publication is based on observations with the MegaPrime/MegaCam, a joint project of the CFHT and CEA/DAPNIA, at the Canada-France-HAwaii Telescope (CFHT), which is operated by the National Research Council of Canada, the institut National des Sciences de l'Univers of the Centre National de la Recherche Scientifique (CNRS), and the University of Hawaii.} In recent years, it has been increasingly recognised that the outskirts of galaxies hold fundamental clues about their formation history. It is into these regions that new material continues to arrive as part of their assembly, by accretion of minor satellites, predominantly at early epochs when large disk galaxies were assembling, as predicted by the currently favored hierarchical formation models. It was also in the outer regions of galaxies that material was deposited during the violent interations in the galaxy's past. Most present-day disk galaxies are suspected to have experienced mergers during the last few billions of years \\citep[e.g.][and references therein]{hammer07}. Based on the systematic deviation of the Milky Way from a number of galaxy scaling relations, \\citet{hammer07} have agrued that our galaxy had most likely escaped any significant major merger event over the last $\\sim 10$ Gyr. These authors suspect that the observed differences between the Milky Way and its neighbour M31 are likely due to the quiescent formation in the former case and to the merger-dominated history for the latter. The observed properties of the stellar content in the outskirts of M31 can be accounted for by either a succession of minor mergers or a major merger, with this material most likely accreting in the most recent half of the age of the Universe \\citep{ibata05}. By analysing the characteristics of spiral galaxy stellar halos formed within a large grid of numerical chemo-dynamical simulations, \\citet{renda05} have shown that at any given total galactic mass, the metallicities of simulated stellar halos span a range in excess of $\\sim$~1~dex. The underlying driver of this metallicity spread can be traced back to the diversity of galactic mass assembly histories. Galaxies with a more extended merging history possess halos which have younger and more metal rich stellar populations than the stellar halos associated with galaxies with a more abbreviated assembly. For a given total mass, galaxies with more extended assembly histories also possess more massive stellar halos. The studies of the Galaxy, and to lesser extent the other large spiral in the Local Group, M31, have been delivering the bulk of the observational constraints on the properties of the stellar content of the outer regions of galaxies. Evidence indicates that the Galaxy might be unrepresentative of a typical spiral galaxy, and it may not even follow the standard scenario of disk formation. The Milky Way halo seems to be populated by old, metal-poor stars, while a few fields in the halo of M31 show a large population of intermediate age stars with a much higher overall metallicity \\citep{brown06}. The fields in M31 in which the above results were obtained have been found to be significantly contaminated by various accretion events \\citep{ibata07} casting doubt on the conclusion that the M31 halo is globally younger and more metal-rich than that of the Milky Way. The current observational evidence therefore demonstrates that halos are complex structures. To establish comprehensively properties of stars in the outskirts of galaxies, and to fully understand their nature and origin, we need to undertake panoramic studies of the outer regions of spiral galaxies beyond the Local Group. To do so, we have obtained deep and wide-field optical imaging data of the nearby early-type spiral M81 (NGC~3031), resolving stars well below the tip of the red giant branch of metal-poor stellar populations. The M81 group of galaxies is one of the nearest groups to our own. It contains one large spiral, two peculiar galaxies (M82 and NGC~3077), two small spirals galaxies (NGC~2976 and IC~2574), as well as a large number of dwarf galaxies \\citep[e.g.][]{karachentsev85, karachentsev01, chiboucas08}. The core galaxies of the group are strongly interacting. Atomic hydrogen observations have revealed the presence of a large number of tidal streams with large, dynamically complex atomic hydrogen clouds embedding M81, M82, NGC~3077, and NGC~2976 \\citep{vdh79,appleton81,yun94,boyce01}. Close interactions between galaxies are capable of leaving tidal debris that could be converted into new stellar systems \\citep[e.g.][]{TT72}. Compared to the Local Group, an interesting feature of the M81 group is the presence of a population of stellar systems dominated by young stars \\citep[e.g.][]{durrell04,demello08a,davidge08}, which are suspected to be of tidal origin \\citep[e.g][]{makarova02}, and which have no counterparts in the Local Group. These young stellar systems, e.g. Holmberg IX, BK 3N, and Garland, are embedded in H{\\sc i} clouds \\citep{boyce01}. Here, we take advantage of our deep and wide field survey to study the spatial distribution of young stellar populations in the outer regions of M81. We report the discovery of new stellar systems in the tidal debris. Analysis of the spatial distribution of old stellar populations, the bi-dimensional distribution, the search for substructures, the metallicity distribution functions, and globular cluster properties over the surveyed area will be reported in forthcoming papers. The stellar populations of immediate interest to the present paper are revealed by the upper main sequence and the red supergiant stars. The layout of this paper is as follows: in Section \\ref{data} briefly represents the data set, while section \\ref{results} studies the young stellar content around M~81. ", "conclusions": "\\label{discussion} Numerous multi-wavelength data sets have unambiguously identified localised regions with signs of current and/or recent star formation activity distributed along H{\\sc i} tidal tails in interacting systems \\citep[e.g.][]{weilbacher03,hibbard05,demello08b}, with a number of these regions suspected to be bounded, the so-called tidal dwarf galaxies \\citep[e.g.][]{duc98,braine01,hancock09}. The star formation activity along the tidal tails appears to be distributed with a similar morphology to H{\\sc i} \\citep[e.g.][]{hibbard05,neff05,hancock07}, in agreement with our finding of the new stellar clumps reported here tracing the dense regions along the southern H{\\sc i} tidal tail. \\citet{hibbard05} have found that UV colours of localised regions of star formation along the tidal tails of the archetypal merging system NGC 4038/39 are consistent with continuing star formation. \\citet{neff05} have argued however, for the case of NGC 7769/71, NGC 5713/19, and the NGC 520 system, that most of young stars in the tails have most likely formed in single bursts. The CMDs of stellar clumps in the debris field of M81, e.g. Arp's loop region, Holmberg IX, contain stars with photometric properties consistent with a wide range of ages, i.e., from $\\sim10\\,$Myr to $\\sim1\\,$Gyr \\citep[e.g.][]{demello08b,sabbi08}, suggesting extended star formation histories. Deep imaging data show however a lack of any concentration of old stars associated with the blue stars, suggesting that the ``old'' stellar component seen in the CMDs of those stellar clumps were formed in the stellar disks of M81 and ejected into the intergalactic medium during tidal passages, whereas the young stars have formed in the tidal debris \\citep{demello08b, weisz08}. The spatial distribution of red giant branch stars over the region connecting M81 and NGC~3077 does not show any noticeable concentrations of old stars associated to the young stellar clumps identified along the southern H{\\sc i} tidal tail. As for, e.g. Holmborg IX and Arp's loop region, this suggests that the old stars seen in the CMDs of the newly reported stellar clumps should have come from one of the interacting systems while, since the stellar clumps along the southern tidal tail are considerably younger than its dynamical age, e.g. $\\sim 250\\,$Myr \\citep{yun99}, young stars formed on site. Fitting single stellar population synthesis models to UV/optical colours of UV-bright stellar substructures within the tidal tails of four ongoing galaxy mergers, \\citet{neff05} found that the star formation appears to be older near the parent galaxies and younger at increasing distances. They have suggested that this could be because the star formation occurs progressively along the tails, or because the star formation has been inhibited near the galaxy/tail interface. \\citet{hibbard05} had reported negative UV and optical colour gradients along the tidal tails of the ``Antennae'' system, indicative of negative age gradients when moving outward along the tails \\citep[see also][]{hibbard01}. Note that the observed colour gradients could be accounted for alternatively by the presence of composite stellar populations. The star formation within the tidal tails in M81 debris field appears to be different. The CMDs of the bulk of stellar clumps in the debris field of M81 indicate that they have ceased forming stars at similar epochs in the past, i.e., $\\sim40$\\,Myr ago. This suggests that the star formation throughout M81 tidal debris field could have been triggered by common events \\citep[see][for a similar conclusion]{davidge08}, and that the physical conditions within the dense regions along the southern H{\\sc i} tidal tail are comparable. The systems within the debris field with younger stars, i.e., Holmberg IX and Garland, are both in the close vicinity of M81 and NGC~3077 respectively, in contrast with the findings of \\citet{neff05}. The diversity of these star formation histories could be most likely related to different gas contents and conditions within the tidal tails of those interacting systems. The exact nature of the previously known young stellar systems within the tidal field of the M81 group, i.e., Holmberg IX, Gerland, and BK 3N, is not entirely clear yet. Detailed modelling of the dynamics of the M81 group suggests that the three largest galaxies in the system had an interaction $\\sim 250$ Myr ago for M81 and NGC~3077, and $\\sim 200$ Myr ago for M82 and M81 \\citep{yun99}. The modelling of optical CMDs of Holmberg IX, BK 3N, and Arp's loop, of comparable depth to ones presented here, suggested that these galaxies have experienced star formation between about 20 and 200 Myr ago \\citep{makarova02}. It has been argued then that these galaxies are tidal dwarf condidates that formed from dust and gas that was blown away from M81 and/or other galaxies in the group \\citep{boyce01,makarova02}. BK3N may be alternatively a pre-existing dwarf irregular galaxy undergoing an interaction with M81 \\citep{boyce01}. The old stellar population associated spatially to Holmberg IX is suspected to belong quite likely to the outer regions of M81, suggesting that this stellar system is of a tidal origin \\citep{sabbi08}. The ages and metallicities of the isochrones that best reproduce the CMDs of the stellar systems along the southern H{\\sc i} arm are similar to these needed to account for the properties of the stellar contents of previously known tidal dwarf candidates in the tidal debris field. This suggests that they both could share similar star formation histories and might be associated to similar events. These isochrone metallicities are significantly higher than the typical metallicities of dwarf galaxies of comparable luminosities \\citep[e.g.][]{rm95}. This suggests that these stellar systems have been assembled from pre-enriched material. This is consistent with the conclusions of \\citet{boone05} who found that the abundances and physical conditions of the molecular complex situated near the line of sight toward Holmberg IX are similar to those found in the disks of spiral galaxies. The distribution of the stellar clumps along an H{\\sc i} tail, tracing the densest clouds within the gaseous arm, indicates that these systems may have been assembled out of gas pulled from one of the large interacting galaxies in the group. A primary criterion for the determination of the nature of these objects is to measure their mass-to-light ratio, which tend to be low for tidal dwarf galaxies due to the absence of dark matter \\citep[e.g.][]{BH92, duc00}. Unfortunately this cannot be measured from the dataset in hand. Without this measurement, we can only conclude that the newly reported stellar clumps are likely (among the nearest) tidal dwarf galaxies. \\begin{table} \\caption{Coordinates of the new young stellar clumps along the southern H{\\sc i} arm. } \\label{gal_prop_obs1} \\begin{tabular}{lll} \\hline ID & R.A. (J2000) & Dec. (J2000) \\\\ \\hline Clump I & 09:57:21.2 & 68:42:55 \\\\ Clump II & 09:59:40.4 & 68:39:19\\\\ Clump III & 10:00:40.4 & 68:39:37 \\\\ \\end{tabular} \\end{table}" }, "1002/1002.2948_arXiv.txt": { "abstract": "Dome A, the highest plateau in Antarctica, is being developed as a site for an astronomical observatory. The planned telescopes and instrumentation and the unique site characteristics are conducive toward Type Ia supernova surveys for cosmology. A self-contained search and survey over five years can yield a spectro-photometric time series of $\\sim$1000 $z<0.08$ supernovae. These can serve to anchor the Hubble diagram and quantify the relationship between luminosities and heterogeneities within the Type Ia supernova class, reducing systematics. Larger aperture ($\\gtrsim$4-m) telescopes are capable of discovering supernovae shortly after explosion out to $z \\sim 3$. These can be fed to space telescopes, and can isolate systematics and extend the redshift range over which we measure the expansion history of the universe. ", "introduction": "Dome A, the highest plateau in Antarctica, is being considered as a site for an optical-to-infrared observatory \\cite{2009astro2010S.308W,2009PASP..121..976S}. There are several attributes that make Dome A attractive for Type Ia supernova (SN Ia) observations compared to temperate sites. The boundary layer is at $\\sim$20 m, above which there is a median optical seeing of 0.3\" (infrared seeing of 0.2''). During the long winter night, there are no disruptions over 24 hours due to weather. The atmospheric column density is relatively low. The sky surface brightness in $K_{\\rm dark} $, from 2.27--2.45 $\\mu$m is expected to be at $\\sim$100 $\\mu\\mbox{Jy arcsec}^{-2}$; OH emission drops out in these wavelengths while sky and telescope thermal emission is suppressed compared to temperate sites as is seen at the South Pole \\citep{1999ApJ...527.1009P} and Dome C \\citep{2005PASA...22..199B}. Dome A has its disadvantages as well. The long day interrupts the longer-term observations necessary for tracking supernova flux evolution. The polar latitude limits the area of sky accessible by a telescope. Auroral activity produces high UV to $B$ band sky emission with non-uniform spatial structure. There are technical hurdles as well, including the remoteness of the site, power, and the transfer of data out of the site, avoiding mirror condensation, and achieving the excellent seeing with the observatory structures. China is moving forward with developing the site \\citep{2009PASP..121..174Y}. In the (austral) summer of 2008-2009 a summer station was constructed. In the next few years, the three 0.5-m Antarctic Schmidt Telescopes (AST3) will be installed. A 1-m class pathfinder is planned to characterize the site and develop the necessary technical knowhow; the intent is to eventually build a 4-10 m-class telescope. In this paper, we explore possible supernova surveys that can be performed by these telescopes. In \\S\\ref{lowz:sec} the search and follow-up of low-redshift supernovae are discussed. In \\S\\ref{highz:sec} the possibilities for a high-redshift search are explored. We summarize", "conclusions": "\\label{conclusions:sec} We have considered how several telescopes planned for Dome A Antarctica can be used to deliver SNe Ia science. The first set of telescopes, AST3 and a 1-m pathfinder, can produce a photometric search and spectro-photometric survey for $z<0.08$ supernovae over an 8000 deg$^2$ region of sky with low airmass and Galactic-dust absorption. Depending on the goals of the survey, an instrument using slitless spectroscopy can provide a spectroscopic time series for all bright transients in the field. However, in this case the fine pixel sampling necessary to minimize sky background and the large survey solid angles lead to enormous numbers of detector pixels. A next-generation large aperture telescope can take advantage of the excellent seeing and the $K_{\\rm dark}$ observing window to discover supernovae out to $z \\sim 3$ and provide spectroscopic typing in a specific redshift window. On-site analysis and telescope control reduce communication needs out of Dome A. However, the more improvement in telemetry, the more that Antarctic advantages can be used in conjunction with transient surveys occurring elsewhere in the world. Dome A discoveries can be promptly announced for observations at other observatories, and external discoveries can be followed at Dome A. More information on site performance would enable better treatment beyond several simplifying assumptions in our calculations. We did not include host-galaxy background nor host-galaxy extinction. Our calculations are based on the fixed median seeing and not the full distribution (which can drop to 0.1\"). We don't simulate the full range of SNe Ia but base our calculations on an average supernova. Exposure time estimates are based on $z=0.08$, the high end of the targeted redshift range. We do not offer precise numbers on how far into twilight supernova observations are possible, i.e.\\ the number of hours per night (we assume 16 hours) and the fraction of the year (we assume five months) that can be spent on the survey. Patterns of suspended observations due to weather are not considered. As more information on the site is collected, we can add more realism into our projections. Nevertheless, these first calculations show that interesting SN~Ia science can be done at Dome A. A number of other science drivers are possible as well. Mapping the peculiar velocity field at $z<0.03$ using SNe Ia is an interesting probe of cosmology \\cite{1997ApJ...488L...1R,2007arXiv0705.0368W}. With a SNe Ia rate of $\\sim 0.006$ $\\mbox{deg}^{-2}\\mbox{yr}^{-1}$, around ten years are required to build significant statistics in the restricted field available at Dome A. Core-collapse supernovae, in particular Type IIP's, are of interest for cosmology \\cite{2006ApJ...645..841N,2009ApJ...694.1067P,2009ApJ...696.1176J}. Exposure times can be tuned to discover these fainter objects and the continuous observing allows monitoring for the UV shock breakout \\cite{2007ApJ...666.1093Q}. The late-time spectroscopy needed to determine velocities of the ejecta to standardize them as distance indicators does limit the time window that these supernovae can be discovered and followed at Dome A. We have not considered the possibility of using adaptive optics at Dome A, although the greater isoplanatic angle and coherence time make this of interest. Such a capability can aid in supernova typing, but it is uncertain how it can achieve precision photometry of point sources on top of an underlying host galaxy over a large field of view. Finally, the interest in an instrument that provides simultaneous optical IFU spectroscopy and NIR imaging of the same field is not unique to the site. Such a camera lessens the need for multi-telescope coordination in building a large set of pan-chromatic supernova data. There remain considerable areas of interest to explore in Antarctic astronomy, supernova surveys, and their intersection. This article has presented a first look at some of the prospects and issues. With the ongoing development at Dome A and the near-term installation of wide-field and 1-m class telescopes, these topics are worth pursuing further." }, "1002/1002.2386_arXiv.txt": { "abstract": "We present a numerical code for calculating the local gravitational self-force acting on a pointlike particle in a generic (bound) geodesic orbit around a Schwarzschild black hole. The calculation is carried out in the Lorenz gauge: For a given geodesic orbit, we decompose the Lorenz-gauge metric perturbation equations (sourced by the delta-function particle) into tensorial harmonics, and solve for each harmonic using numerical evolution in the time domain (in 1+1 dimensions). The physical self-force along the orbit is then obtained via mode-sum regularization. The total self-force contains a dissipative piece as well as a conservative piece, and we describe a simple method for disentangling these two pieces in a time-domain framework. The dissipative component is responsible for the loss of orbital energy and angular momentum through gravitational radiation; as a test of our code we demonstrate that the work done by the dissipative component of the computed force is precisely balanced by the asymptotic fluxes of energy and angular momentum, which we extract independently from the wave-zone numerical solutions. The conservative piece of the self-force does not affect the time-averaged rate of energy and angular-momentum loss, but it influences the evolution of the orbital phases; this piece is calculated here for the first time in eccentric strong-field orbits. As a first concrete application of our code we recently reported the value of the shift in the location and frequency of the innermost stable circular orbit due to the conservative self-force [Phys.\\ Rev.\\ Lett.\\ {\\bf 102}, 191101 (2009)]. Here we provide full details of this analysis, and discuss future applications. ", "introduction": "The prospects for detecting gravitational waves from the inspiral of compact objects into massive black holes have motivated, over the past decade, research in effort to understand the general-relativistic orbital evolution in such systems. The underlying elementary theoretical problem is that of a pointlike mass particle in a strong-field orbit around a Kerr black hole of a much larger mass. The dynamics of such systems can be described in a perturbative fashion in terms of an effective gravitational self-force (SF) \\cite{Mino:1996nk,Quinn:1996am,Poisson:2003nc,Gralla:2008fg,Pound:2009sm}; knowledge of this force is a prerequisite for describing the precise evolution of the orbit and the phasing of the emitted gravitational waves. There is an active research program focused on the development of computational methods and actual working codes for the SF in Kerr spacetime \\cite{Barack:2009ux}. This research agenda is being pursued in incremental steps, through exploration of a set of simplified model problems with increasing complexity and physical relevance. Much of the initial work has concentrated on a scalar-field toy model \\cite{Burko:2000xx,Barack:2000zq,Detweiler:2002gi, DiazRivera:2004ik,Barack:2007jh,Haas:2007kz,Vega:2007mc}, but more recently workers have begun to tackle the gravitational case \\cite{Barack:2002ku,Barack:2007tm,Keidl:2006wk,Detweiler:2008ft,Berndtson:2009hp}. The state of the art is represented by three independent calculations of the gravitational SF for circular geodesic orbits in Schwarzschild geometry \\cite{Barack:2007tm,Detweiler:2008ft,Berndtson:2009hp}. These calculations use different analytic and numerical methods (and they even invoke different physical interpretations of the SF), but they were shown to be fully consistent with each other \\cite{Sago:2008id,Berndtson:2009hp}. These calculations were also shown to be consistent with results from post-Newtonian theory in the weak-field limit \\cite{Detweiler:2008ft,Blanchet:2009sd,Damour:2009sm}. In the current work we extend the analysis of Ref.\\ \\cite{Barack:2007tm} (hereafter ``Paper I'') from the special class of circular geodesics to generic (bound) geodesics of the Schwarzschild geometry. This generalization is astrophysically relevant because real inspirals often remain quite eccentric up until the eventual plunge into the massive hole \\cite{Barack:2003fp}. At a more fundamental level, the generalization to eccentric orbits is interesting because it allows us to start exploring in earnest the conservative effects of the SF---for instance, how it influences the orbital precession. Eccentric orbits have already been considered in calculations of the scalar \\cite{Haas:2007kz} and electromagnetic (EM) \\cite{Haas:Capra} SFs by Haas. While these calculations are of a less direct astrophysical relevance, they offer an important test-ground for computational techniques potentially applicable in the gravitational problem too. Indeed, many elements of our numerical method take their inspiration from Haas' work. The numerical code we present here takes as input the two orbital parameters of an eccentric Schwarzschild geodesic (the semi-latus rectum and eccentricity, to be defined below), and returns the value of the Lorenz-gauge gravitational SF along this geodesic. The dissipative and conservative pieces of the SF are returned separately. Here we do {\\em not} consider the evolution of the orbit under the effect of the SF, but leave this important next step for future work. We envisage using, to this end, a version of the ``osculating geodesics'' method \\cite{Pound:2007th}, which takes as input the value of the SF along geodesics tangent to the actual inspiral orbit. A systematic framework for analyzing the long-term evolution of the inspiral orbits, using multiple-scale perturbation methods, was recently developed by Hinderer and Flanagan \\cite{Hinderer:2008dm} (cf.~Sec.~VII of Gralla and Wald \\cite{Gralla:2008fg}). Our strategy is similar to that of Paper I. Its basic elements are (i) the Lorenz-gauge perturbation formalism of Barack and Lousto \\cite{Barack:2005nr}, (ii) a finite-difference algorithm for numerical integration of the Lorenz-gauge perturbation equations in the time domain, and (iii) mode-sum regularization \\cite{Barack:1999wf,Barack:2001bw,Barack:2001gx,Barack:2002bt}. The perturbation formalism is based on a tensor-harmonic decomposition of the perturbed Einstein equations in the Lorenz gauge. The equations are augmented with ``gauge damping'' terms designed to suppress gauge violations \\cite{Barack:2005nr}, and are written as a set of 10 hyperbolic equations (for certain linear combinations of metric components) which do not couple at their principal parts. These equations are sourced by the (tensor-harmonic modes of the) particle's energy-momentum, modeled with a delta-function distribution along the specified eccentric geodesic. The equations are solved numerically mode by mode in the time domain using characteristic coordinates on a uniform 1+1-dimensions mesh. The non-radiative monopole and dipole modes cannot be evolved stably in this manner; instead, we solve for these two modes separately in the frequency domain, using the recently introduced ``extended homogeneous solutions'' technique \\cite{Barack:2008ms} to cure the irregularity of the Fourier sum near the particle. The code records the value of the perturbation modes and their derivatives along the orbit (each mode has a $C^0$ behavior at the particle and hence a well-defined value there, as well as a well-defined ``one-sided'' derivatives). These values are then fed into the ``mode-sum formula'' \\cite{Barack:2001gx}, which returns the physical SF through mode-by-mode regularization. One of the primary advantages of the time-domain approach is that eccentric orbits---even ones with large eccentricity---are essentially ``as easy'' to deal with as circular orbits, with computational cost being only a weak function of the eccentricity \\cite{Barton:2008eb}. Also, a time-domain code for circular orbits can be upgraded with relative ease to accommodate eccentric orbits (such a generalization is radically less straightforward in the frequency domain). Still, there are several important technical issues which arise in the time-domain upgrade from circular to eccentric orbits, and need to be addressed. We list some of these issues below. \\begin{itemize} \\item Most obvious, the computational burden increases significantly because the parameter space for geodesics turns from 1D (circular) to 2D (eccentric). Moreover, for each given geodesic parameters the SF becomes a function along the orbit (it has a constant value along a circular geodesic), and one is required to obtain this function over an entire radial period. The latter becomes a technical hurdle in situations where the radial period is very large---e.g., close to the last stable orbit, or for orbits with very large radii. \\item In paper I we were able to improve the convergence rate of our finite-difference algorithm using a Richardson-type extrapolation to the limit of a vanishing numerical grid-cell size. That was possible because in the circular-orbit case the numerical mesh could be easily arranged such that the local discretization error varied smoothly along the orbit. This cannot be achieved in any simple way when the orbit is eccentric, and as a result one cannot implement a similar Richardson extrapolation. The practical upshot is that one is forced to implement a higher-order finite-difference scheme: a 2nd-order-convergent algorithm (as in Paper I) proves insufficient in practice. For this work we developed an algorithm with a 4th-order global convergence. The algorithm takes a rather complicated form near the particle's trajectory, where the field (the Lorenz-gauge metric perturbation) has discontinuous derivatives. To somewhat lessen this complexity (and reduce the number of grid points needed as input for the finite-difference formula) the algorithm makes use of suitable junction conditions across the orbit. The eventual numerical scheme is considerably more sophisticated---and involved---compared to that of Paper I. \\item In the mode-sum scheme one first calculates the contribution to the ``full'' (pre-regularization) force from each tensorial-harmonic mode of the perturbation, and then decomposes this into {\\it spherical} harmonics. The necessary input data for the mode-sum formula are the individual spherical-harmonic contributions. This procedure involves the implementation of a tensor--scalar coupling formula, whose details depend on the orbit in question. The coupling formula simplifies considerably in the circular-orbit case; it reverts to its full complicated form [Eq.\\ (\\ref{eq:Flfull}) with Appendix \\ref{app:Ffull}] when eccentric orbits are considered. \\item The computation of the monopole and dipole contributions to the SF (which we perform in the frequency domain, as mentioned above) becomes much more involved in the eccentric-orbit case. First, the spectrum of the orbital motion now includes all harmonics of the radial frequency, and one has to calculate and add up sufficiently many of these harmonics. A second, more technically challenging complication arises from the fact that the perturbation becomes a non-smooth function of time across the orbit (at a given radius), which disrupts the high-frequency convergence of the Fourier sum at the particle (a behavior reminiscent of the Gibbs phenomenon). A general method for circumventing this problem in frequency-domain calculations was devised recently in Ref.\\ \\cite{Barack:2008ms}, and we implement it here for the first time. \\item In exploring the physical consequences of the SF it is useful to split the SF into its dissipative and conservative pieces, and discuss their corresponding effects in separate. This splitting is straightforward in the circular-orbit case: The conservative piece is precisely the (Schwarzschild) $r$ component of the SF, while the (Schwarzschild) $t,\\varphi$ components exactly account for the entire dissipative effect. This is no longer true for eccentric orbits, where each of the Schwarzschild components mixes up both dissipative and conservative pieces, and it is not immediately obvious how to extract these pieces individually. Here we suggest and implement a simple new method for constructing the dissipative and conservative pieces out of the computed Schwarzschild components of the SF (without resorting to a calculation of the advanced perturbation). The method takes advantage of the general symmetries of Schwarzschild geodesics. \\end{itemize} With the computational framework in place, we can start to explore the physical effects of the gravitational SF. In this article we concentrate on two such effects. First, we calculate the loss of orbital energy and angular momentum, over one radial period, due to the dissipative piece of the SF. We extract these quantities directly from the computed SF along the geodesic orbit (for a sample of orbital parameters). These ``lost'' energy and angular momentum must be balanced by the total amount of energy and angular momentum in the gravitational waves radiated to spatial infinity and into the black hole over a radial period. We derive formulas for extracting these quantities from the far-zone and near-horizon numerical Lorenz-gauge solutions, and demonstrate numerically that they agree well with the values computed from the local SF. Our values for the energy and angular momentum losses also agree with those previously obtained by others using other methods. The second effect we consider is conservative, and cannot be inferred indirectly from the asymptotic gravitational waves: It is the conservative shift in the location and frequency of the Innermost Stable Circular Orbit (ISCO). The analysis of the ISCO shift requires knowledge of the SF along slightly eccentric geodesics near the last stable orbit, and our code provides the necessary SF data for the first time. We reported the results in a recent Letter \\cite{Barack:2009ey}, and here we describe our analysis in full detail. The quantitative determination of the ISCO shift is important in that it provides a strong-field benchmark for calibration of approximate (e.g., post-Newtonian) descriptions of binary inspirals. Our result for the ISCO frequency shift has already been incorporated by Lousto {\\it et al.}~in their ``empirical'' fitting formula for predicting the remnant mass and spin parameters in binary mergers \\cite{Lousto:2009mf,Lousto:2009ka}; and by Damour \\cite{Damour:2009sm} for breaking the degeneracy between certain unknown parameters of the Effective One Body (EOB) formalism. Perhaps of a more direct relevance to the problem of the phase evolution in binaries with extreme mass-ratio is the effect of the SF on the periapsis precession of the eccentric orbit---also a conservative effect. SF corrections to the precession rate have been analyzed for weak-field orbits and within the toy model of the EM SF \\cite{Pound:2005fs,Pound:2007ti}, but never before for the gravitational problem in strong field. Our code generates the SF data necessary to tackle this problem for the first time. We leave the detailed analysis of SF precession effects to a forthcoming paper. The paper is organized as follows. In Sec.\\ II we review the relevant theoretical background: bound geodesics in Schwarzschild geometry, the Lorenz-gauge metric perturbation formulation, and the construction of the SF via the mode-sum formula. Section III describes our numerical method in detail, and in Sec.\\ IV we present numerical results for a few sample eccentric orbits, including a ``zoom--whirl'' orbit. We explain how the dissipative and conservative pieces of the computed SF can be extracted from the numerical data, and present these two pieces separately in a few sample cases. We also analyze the dissipative effect of the SF and demonstrate the consistency between the dissipated energy and angular momentum inferred from the local SF, and that extracted from the asymptotic gravitational waves. Section V covers the ISCO-shift analysis, and in Sec.\\ VI we summarize and discuss future applications of our code. Throughout this work we use standard geometrized units (with $c=G=1$), metric signature ${-}{+}{+}{+}$, and (unless indicated otherwise) Schwarzschild coordinates $x^\\mu = (t,r,\\theta,\\varphi)$. ", "conclusions": "This work marks a new frontline in the program to model realistic two-body inspirals in the extreme mass-ratio regime. For the first time we are able to calculate the full [$O(\\mu^2)$] gravitational SF across (essentially) the entire parameter space of strong-field bound geodesics in Schwarzschild spacetime. This work also represents a first complete end-to-end implementation of a range of computational techniques which were developed gradually over the past decade: mode-sum regularization scheme \\cite{Barack:2001gx}, the 1+1D Lorenz-gauge perturbation formalism \\cite{Barack:2005nr}, and the method of extended homogeneous solutions \\cite{Barack:2008ms}. As the reader may appreciate, the underlying computational challenge is rather daunting, given the complexity of the field equations, the technical subtleties involved in dealing with the delta-function source, the high computational cost, and the need to patch together different techniques in both the time and frequency domains. Our eventual working code is of considerable complexity and took over two years to develop and test. Following is a summary of the various tests which helped us establish confidence in our code's performance. (i) The mode-sum regularization procedure is self-validating, in the sense that it is extremely sensitive to errors in the computation of the perturbation multipoles (especially the high-$l$ ones, which are most computationally demanding). If the regularized mode sum shows the expected fall-off behavior at large $l$, this by itself is a strong indication that the high-$l$ modes were calculated correctly. (ii) The code reproduces the known results in the circular-orbit case; these results are now confirmed by 3 independent analyses \\cite{Barack:2007tm,Detweiler:2008ft,Berndtson:2009hp}. (iii) Our evolution code reproduces the correct asymptotic fluxes of gravitational-wave energy and angular momentum, as verified by comparing with results in the literature. (iv) The work done by the dissipative piece of the computed SF is found to precisely balance these fluxes. (v) The value of the ISCO frequency shift derived from our SF seems consistent with the value derived in EOB at 3rd post-Newtonian (PN) order: Damour recently showed that the latter is about 72.5\\% of the SF value, with the difference likely attributed to higher-order PN terms \\cite{Damour:2009sm}. In principle, our code can return the SF along any bound geodesic in Schwarzschild geometry, although in practice computational cost may becomes prohibitive when the orbital period is too large (i.e., for very large $p$ and/or $e$ close to unity). The ``workable domain'' of our code using a current-day standard single-processor desktop computer is roughly $0\\lesssim e\\lesssim 0.5$ and $p\\lesssim 20M$ if a fractional accuracy of $<10^{-4}$ in the SF is sought. The current algorithm incorporates an explicit reference to the radial frequency parameter, so it cannot be used to tackle unbound orbits. However, it may be adapted with moderate effort to handle unbound orbits (including orbits below the last stable orbit) as well. Even within the above ``workable domain'', the current code is discouragingly slow. It takes a few days to compute the SF along a single strong-field geodesic, which makes it impractical to cover the entire parameter space of inspirals at sufficient resolution (having in mind the development of theoretical gravitational waveform templates for astrophysical inspirals). There are, however, various ways in which one may improve the computational efficiency and speed. Most obvious, one can use distributed computing---our algorithm should be easily amenable for distribution on a cluster, since different $l$ modes can be calculated in parallel. Other natural approaches include the use of mesh refinement (see Thornburg's recent report \\cite{Thornburg:2009mw}) and/or higher-order finite-difference schemes. One may also seek to reduce the computational cost attached to the initial stage of the numerical evolution (when spurious initial waves dominate) by iteratively improving the initial conditions for the evolution---this idea is already being implemented successfully in a 2+1D framework \\cite{BDprep}. Other ideas represent a more significant deviation from our approach: (i) work entirely in the frequency domain, making full use of the method of extended homogeneous solutions \\cite{Barack:2008ms} (this is likely to prove most efficient with smaller eccentricities); or (ii) abandon finite differencing altogether and instead use finite elements or other pseudospectral techniques \\cite{Canizares:2009ay,Field:2009kk}, benefiting from their natural flexibility in accommodating multiple lengthscales. Nonetheless, some interesting applications are already possible with the current version of our code, and we have presented one of them in Sec.\\ \\ref{Sec:ISCO}. Our computation of the ISCO frequency shift represents the first physically-meaningful new result coming out of the SF program, and it has already informed both Numerical-Relativistic calculations \\cite{Lousto:2009mf,Lousto:2009ka} and EOB/PN theory \\cite{Damour:2009sm}. Further ideas for SF/EOB synergy were recently discussed by Damour in \\cite{Damour:2009sm}. In particular, Damour showed that a computation of the (gauge invariant) conservative SF correction to the precession rate of the periapsis, for slightly eccentric orbits, will give access to the presently unknown 4PN (and possibly higher-order) parameters of EOB theory. We are currently working to extract the necessary SF data to facilitate this calculation [these are essentially the coefficients $F_0^r$, $F^r_{1}$ and $F^1_{\\varphi}$ of Eqs.\\ (\\ref{eq160}) and (\\ref{eq150}), as functions of $r_0$]. Our preliminary results show excellent agreement with the analytic EOB predictions at 2PN and 3PN, and we are hoping to publish these findings elsewhere \\cite{EOBGSF}. More generally, our code allows to tackle the calculation of post-geodesic [$O(\\mu)$] precession effects at {\\em any} eccentricity, not necessarily small. This would not only provide a handle on the `Q' function of EOB theory, but can also directly inform the computation of conservative effects in inspiral trajectories. Information on the periastron advance as a function of $p,e$ could, for example, be incorporated into an orbital evolution scheme \\`{a} la Pound \\& Poisson \\cite{Pound:2007th}. We are currently investigating this direction. Finally, we comment briefly on the possibility of extending this work to the Kerr case (a more elaborate discussion of possible strategies for attacking the Kerr problem can be found in \\cite{Barack:2009ux}). The framework of the current code, i.e., numerical evolution in 1+1D is not directly applicable in Kerr spacetime, because the perturbation equations in Kerr cannot be separated into harmonics in the time domain. It is possible to tackle the field equations in 2+1D (i.e, separating the perturbation into azimuthal $m$-modes only) or in full 3+1D, and there has been considerable progress in that direction in the past two years---although work so far has been restricted to the toy problem of a scalar field in Schwarzschild geometry \\cite{Barack:2007jh,Lousto:2008mb,Vega:2009qb,BDprep}. Schemes for regularizing the SF directly in 2+1D or 3+1D have been proposed and recently implemented \\cite{Barack:2007we,Vega:2009qb,BDprep}. Alternatively, one may attempt to tackle the field equations in 1+1D by properly accounting for the coupling between different $l$-harmonics, and a similar strategy may be applicable in the frequency domain too. A first calculation of the SF in the Kerr case (using the frequency-domain approach)---for a scalar charge in a circular equatorial orbit---will be presented in a forthcoming paper \\cite{WBprep}." }, "1002/1002.0510_arXiv.txt": { "abstract": "We show that quintessence, when it is described by a tachyonic field, can amplify a tiny primordial gradient generating a preferred direction in the sky. In its simplest realization, this mechanism only affects the Cosmic Microwave Background fluctuations at the quadrupole level. We briefly discuss how higher multipoles can also be affected, once the full structure of the quintessence potential is taken into account. ", "introduction": "% The formulation of the cosmological principle~\\cite{Einstein:1917ce} coincides with the birth of modern scientific cosmology. Over the last decades, homogeneity and isotropy of the Universe have been tested to increasing degrees of accuracy. In particular, the radiation of the Cosmic Microwave Background (CMB) is one of the most sensitive probes to test the isotropy of the Universe. The observed temperature of the CMB is observed to be uniform at zeroth order (once the dipole is subtracted): this result led to the formulation of the theory of primordial inflation. Moreover, approximate statistical isotropy appears to hold even for the small fluctuations of the CMB temperature. In spite of these striking results, several analyses of the CMB fluctuation maps, starting from~\\cite{Eriksen:2003db,Hansen:2004vq} (see, {\\em e.g.}, \\cite{Hanson:2009gu} for a more recent analysis), have shown the existence of anomalies associated to some degree of breaking of statistical isotropy. Even though a clear consensus on the subject is still absent ({\\em e.g.}, see \\cite{Bennett:2010jb} for an up to date discussion) intensive work has been done in the past also as a response to the theoretical challenge of naturally generating a preferred directions in the sky. The models proposed usually rely either on early-Universe or on late-Universe mechanisms. While the former lead to an {\\em intrinsic} anisotropy in the primordial spectrum of metric perturbations, the latter produce anisotropy in the {\\em observed} CMB via anisotropic Sachs-Wolfe effect. Early-Universe mechanisms include~\\cite{Ackerman:2007nb} (see however~\\cite{Himmetoglu:2008hx}) and~\\cite{Watanabe:2009ct}, that rely on vector fields or~\\cite{Boehmer:2007ut}, that rely on spinor fields, as well as~\\cite{Donoghue:2007ze,Erickcek:2008sm}, that rely on primordial gradients. Late-Universe mechanisms can invoke magnetic fields~\\cite{Campanelli:2007qn} or anisotropies in the dark energy equation of state~\\cite{Rodrigues:2007ny,Koivisto:2008ig,Battye:2009ze}. Despite those many attempts, it is fair to say that breaking statistical isotropy usually requires strong and aesthetically unappealing assumptions. In the present work we show a (relatively) natural mechanism that can give rise to a preferred direction in the CMB sky. This mechanism is a hybrid of early and a late Universe ones and does not rely on vector fields or on unusual properties of the dark energy sector, but rather on a simple model of scalar dark energy, characterized by a non-trivial, tachyonic ($V''(\\phi)<0$) potential. Models of tachyonic quintessence frequently appear in the literature. For instance, if the quintessential scalar is described by a pseudo-Nambu-Goldstone Boson (pNGB)~\\cite{Frieman:1995pm}, then half of the extrema of the potential $V(\\phi)\\propto 1+\\cos(\\phi/f)$ are tachyonic. Another scenario of tachyonic quintessence was proposed in~\\cite{Kallosh:2002gf}, where the magnitude of the tachyonic mass is related to the height of the potential at its maximum. Our main assumption is that a tachyonic field $\\phi$ has a small primordial gradient. As the Hubble parameter drops below $\\sqrt{\\left|V''(\\phi)\\right|}$, the initial gradient is amplified by the ``pull'' of its tachyonic potential, effectively converting a small primordial isocurvature perturbation into a larger curvature mode. In first approximation, this leads to an ``ellipsoidal'' Universe~\\cite{Campanelli:2007qn} which could explain the low quadrupole CMB amplitude observed in both COBE and WMAP data~\\cite{Contaldi:2003zv}, if the preferred direction in the quintessence correlated with that of the primordial quadrupole. As we will see, our mechanism will be at work for a tachyonic mass a few times larger than $\\Lambda/M_\\mathrm{Pl}^2$. This is, in particular, the case for a pNGB with value of the axion constant $f$ slightly sub-Planckian. The primordial gradient in $\\phi$ can be a remnant of the epoch that preceded the last bout of inflation, provided that such last bout were sufficiently short (this argument will be discussed in greater detail in section~\\ref{gradient}). In this respect, our scenario could shed some light on pre-inflationary dynamics, similarly to the situations studied, {\\em e.g.}, in~\\cite{Chang:2007eq} and especially~\\cite{Donoghue:2007ze}. In our case, a constant gradient can be the dominant remnant of a more general inhomogeneous initial value of the tachyonic field, {\\em i.e.}, we can expand $\\phi\\approx \\kappa_0+\\kappa_1\\,z_p+\\kappa_2\\,z^2_p+\\dots$ -- for sake of simplicity, we will only consider powers of the $z$ coordinate, though the argument for the dominance of the linear term would still hold, were more generic terms considered -- at $N_\\mathrm{i}\\equiv N_{\\mathrm{obs}}+\\delta N$ e-foldings before the end of inflation ($N_{\\mathrm{obs}}$ corresponds to the time when the current cosmological scales left the horizon and $z_p$ to the physical distance along the $z$ direction). We assume that each term $\\kappa_k\\,z_p^k$ had given an equal contribution to the energy in $\\phi$ at that moment. Today, the scales, that left the horizon $N_{\\mathrm{obs}}+\\delta N$ e-foldings before the end of inflation, are still outside the horizon by a factor of $\\e^{\\delta N}$, {\\em i.e.}, they correspond to physical distances of the order of $\\e^{\\delta N}\\,H_0^{-1}$. Assuming no (other) significant evolution in $\\phi$~\\footnote{The evolution in $\\phi$ occurs only at sub-horizon scales and in the late Universe, when the Hubble parameter is smaller than the (absolute value of the tachyonic) mass of quintessence, and it is not sufficient to invalidate this argument.}, the assumption that each term $\\kappa_k\\,z_p^k$ be of the same order implies that $\\kappa_k\\approx \\e^{-k\\,\\delta N}\\,H_0^k$. As a consequence, inside the horizon $z_p\\approx H_0^{-1}$, each term $\\kappa_k\\,z_p^k$ contributes like $\\e^{-k\\,\\delta N}$ to the energy in $\\phi$. This implies that the terms with the lowest powers in $z_p$ will give the largest contributions to the metric perturbations at sub-horizon scales. For this reason, our analysis will be focused on the approximation where, in the early Universe, $\\phi=\\kappa_1\\,z$ with $\\phi(z=0)$ being set equal to zero by an appropriate choice of the origin of the $z$ coordinate.\\\\ The paper is organized as follows. In section~\\ref{metric} we set up our system and solve the corresponding Einstein equations. Sections~\\ref{rs} and~\\ref{cmb} contain the effect of our anisotropic metric on the redshift and in particular on the CMB quadrupole anisotropy. Section~\\ref{gradient} contains a discussion of the magnitude of the primordial gradient responsible for the anisotropy. Finally, in the concluding section we will discuss how our mechanism could also lead to alignments of the higher multipoles. ", "conclusions": "% If inflation lasted a relatively short time, then some primordial gradients could exist as a relic of the chaotic pre-inflationary dynamics. A gradient in the quintessence field, even if too small to leave any detectable effects during the observable epoch of inflation, might be amplified by a tachyonic quintessence to have observable magnitude today. We have seen that in the simplest scenario, the dominant effect of the amplified gradient is on the CMB quadrupole, whose observed amplitude sets the strongest constraints on the parameters of the problem. The fact that CMB fluctuations are affected only at the quadrupole level is a consequence of the approximation $\\phi\\simeq \\kappa_1(t)\\,z$. Higher powers of $z$ in the expansion of $\\phi$ would lead to the generation of higher multipole contribution. Generically, terms of order $z^n$ in the expansion of the scalar field $\\phi(t,z)$ will generate contributions on both $\\delta a(t,z)$ and $\\delta b(t,z)$ up to order $z^{n-1}$, which, in turn, provide terms up to order $z^{n+1}$ in $\\bar\\zeta_\\mathrm{iSW}$ as for equation~\\eqref{zeta}. That is, terms of order $z^n$ in $\\phi(t,z)$ will affect $C_\\ell$'s of $\\ell=n+1$. Given the hints of multipole alignments up to $\\ell\\sim 40$, it would be interesting to be able to extend our mechanism in such a way that larger powers of $z$ in the expression of $\\phi(t,\\,z)$ are generated. One possibility is that the self-interactions of $\\phi$ during the recent cosmological evolution are responsible. For instance, if the quintessence field $\\phi$ is given by a pseudo-Nambu-Goldstone boson, its potential is $V(\\phi)=\\Lambda\\,\\left[1+\\cos\\left(\\phi/f\\right)\\right]/2$ ($f^2=\\Lambda/(2\\,m^2)$). By following our analysis of section~\\ref{metric}, we see that, because of the Klein-Gordon equation, the coefficients of lower powers in $z$ will act as sources for those of higher powers, hence generating higher multipoles. A further investigation of this mechanism is subject of future work.\\\\ {\\bf Acknowledgments.} This work has been supported in part by the NSF grant PHY - 0855119. We thank John Donoghue, Burak Himmetoglu, David Langlois, Nemanja Kaloper and Marco Peloso for interesting discussions." }, "1002/1002.4619_arXiv.txt": { "abstract": "The new high energy data coming mainly from the {\\it Fermi} and {\\it Swift} satellites and from the ground based Cerenkov telescopes are making possible to study not only the energetics of blazar jets, but also their connection to the associated accretion disks. Furthermore, the black hole mass of the most powerful objects can be constrained through IR-optical emission, originating in the accretion disks. For the first time, we can evaluate jet and accretion powers in units of the Eddington luminosity for a large number of blazars. Firsts results are intriguing. Blazar jets have powers comparable to, and often larger than the luminosity produced by their accretion disk. Blazar jets are produced at all accretion rates (in Eddington units), and their appearance depends if the accretion regime is radiatively efficient or not. The jet power is dominated by the bulk motion of matter, not by the Poynting flux, at least in the jet region where the bulk of the emission is produced, at $\\sim 10^3$ Schwarzschild radii. The mechanism at the origin of relativistic jets must be very efficient, possibly more than accretion, even if accretion must play a crucial role. Black hole masses for the most powerful jets at redshift $\\sim3$ exceed one billion solar masses, flagging the existence of a very large population of heavy black holes at these redshifts. ", "introduction": "With the launch of the {\\it Fermi} satellite we have entered in a new era of blazar research. The Large Area Telescope (LAT) onboard {\\it Fermi} has $\\sim$20 fold better sensitivity than its predecessor EGRET in the 0.1-100 GeV energy range, enabling the detection of hundreds (and thousands in the end) of blazars. For the brightest we can study not only their high state, but also their more normal and quiescent states. Other key information come from the UVOT, XRT and BAT telescopes onboard the {\\it Swift} satellite, covering the optical and X--ray energy range of the blazars detected by {\\it Fermi}. Together with data gathered by ground based telescope, we can routinely assemble simultaneous spectral energy distributions (SEDs). Blazars are extremely variable sources, so having simultaneous snapshot of the entire SED is a mandatory to start to meaningfully explore their physics. Up to now, there exist two catalogues of {\\it Fermi} blazars. The first \\cite{abdo1} lists the $\\sim$100 blazars detected with high significance ($>10\\sigma$) during the first three months of all sky survey ({\\it Fermi} patrols the entire sky every three hours, i.e. two orbits), while the second, just published \\cite{abdo2} list the $\\sim$700 blazars detected at more than 4$\\sigma$ during the first 11 months of operation. In addition, there is another catalogue of blazars detected by the {\\it Swift}/BAT instruments in the 15--55 keV energy range \\cite{ajello2009}, listing 38 blazars at high galactic latitudes ($|b<15^\\circ|$, see also \\cite{cusumano2009} in which blazars at lower Galactic latitudes are present). \\begin{figure} \\includegraphics[height=.35\\textheight]{0836_710f} \\includegraphics[height=.35\\textheight]{gfos_bat_lat} \\caption{Left panel: the SED of the FSRQ 0836+710. In this very powerful source the two broad humps produced by the jet leave the disk component ``naked\" and well visible. The line is a fitting model. Adapted from \\cite{chasing}. Right panel: The blazar sequence \\cite{fossati1998}, \\cite{gg1998}, with overplotted the LAT (0.1--100 GeV) and BAT (15--55 keV) energy ranges. } \\label{sequence} \\end{figure} In the following I will report on the results obtained analysing all blazars with redshift of the 3--months {\\it Fermi} list (85 sources, excluding 4 FSRQs with poor data coverage) and on the 10 FSRQs with $z>2$ present in the BAT list. The novelty, with respect to previous studies, is not only the increased quantity and quality of the data, but also the realisation that, at least for high power objects, it is possible to find the black hole mass and the accretion rate. To understand how this is possible, consider that all blazars have a SED characterised by two broad emission humps, located at smaller frequencies as the bolometric luminosity increases. The first hump is thought to be produced by the synchrotron process, while the high energy one is though to be due to the Inverse Compton process. At low luminosities, the two humps have more or less the same power, while the high energy hump becomes more dominant when increasing the bolometric luminosity (this is the so called {\\it blazar sequence} illustrated in Fig. \\ref{sequence} \\cite{fossati1998}, \\cite{gg1998}). Powerful (broad line emitting) FSRQs have the synchrotron hump at sub--mm frequencies, and the high energy peak in the MeV band. They should thus have {\\it steep} [$\\alpha_\\gamma>1$; $F(\\nu)\\propto \\nu^{-\\alpha}$] spectra in the {\\it Fermi} band, and {\\it flat} spectra ($\\alpha_X<1$) in the BAT band. If the same electrons are also responsible for the low energy hump, then the emission above the synchrotron peak should have a spectrum as steep as observed in the {\\it Fermi} band. As a consequence, the synchrotron IR--optical--UV flux is weak, and the emission from the accretion disk can easily dominate in these bands. The left panel of Fig. \\ref{sequence} illustrates this point by showing the SED of the powerful FSRQ 0836+710. By fitting the IR-opt--UV data with a Shakura--Sunjaev \\cite{shakura1973} disk we can find both the black hole mass and the accretion rate. At the other extreme of the blazar sequence (i.e. low power BL Lacs), there is no sign of thermal disk emission, nor of emission lines produced by the photo--ionising disk photons. For them we can only derive an upper limit of the accretion luminosity. We model the non--thermal SED of all blazars with a one--zone, synchrotron and Inverse Compton leptonic model, to find the physical parameter of the emitting region of the jet, including the power transported in the form of particles and fields \\cite{ggcanonical}. We use a cosmology with $h=\\Omega_\\Lambda=0.7$ and $\\Omega_{\\rm M}=0.3$, and use the notation $Q=Q_X 10^X$ in cgs units (except for the black hole masses, measured in solar mass units). ", "conclusions": "" }, "1002/1002.3380_arXiv.txt": { "abstract": "Oscillons are extremely long-lived, oscillatory, spatially localized field configurations that arise from generic initial conditions in a large number of nonlinear field theories. With an eye towards their cosmological implications, we investigate their properties in an expanding universe. We (1) provide an analytic solution for one-dimensional oscillons (for the models under consideration) and discuss their generalization to three dimensions, (2) discuss their stability against long wavelength perturbations, and (3) estimate the effects of expansion on their shapes and lifetimes. In particular, we discuss a new, extended class of oscillons with surprisingly flat tops. We show that these flat-topped oscillons are more robust against collapse instabilities in (3+1) dimensions than their usual counterparts. Unlike the solutions found in the small amplitude analysis, the width of these configurations is a nonmonotonic function of their amplitudes. ", "introduction": "\\label{Intro} A number of physical phenomenon from water waves traveling in narrow canals \\cite{Russell:1844}, to phase transitions in the early Universe \\cite{Frieman:1988ut} exhibit the formation of localized, energy density configurations, even without gravitational interactions. The reason for their longevity are varied. Some configurations are stable due to conservation of topological or nontopological charges, while some are due to a dynamical balance between the nonlinearities and dissipative forces. Relativistic, scalar field theories (with nonlinear potentials) form simple yet interesting candidates for studying such phenomenon. Some well-studied examples include topological solitons in the $1+1$-dimensional Sine-Gordon model and nontopological solitons such as Q-balls \\cite{Coleman:1985ki}. The Sine-Gordon soliton is stationary in time whereas the Q-balls are oscillatory in nature. Both have conserved charges which make them stable (at least without coupling to gravity). This paper deals with another interesting example of such localized configurations called oscillons (also called breathers). Like the Sine-Gordon soliton, they can exist in real scalar fields, and like the Q-balls they are oscillatory in nature. Unlike both of the above examples they do not have any known conserved charges (however, see \\cite{Kasuya:2002zs} for an adiabatic invariant). In general they decay, however their lifetimes are significantly longer than any natural time scales present in the lagrangian. Along with their longevity, another fascinating aspect of oscillons is that they emerge naturally from relatively arbitrary initial conditions. Not all scalar field theories support oscillons. In the next section we discuss the requirements for the potential. Here, we note that the requirement is satisfied by a large number of physically well-motivated examples. For example, the potential for the axion, as well as almost any potential near a vacuum expectation value related to symmetry breaking, support oscillons. Oscillons have also been found in the restricted standard model $SU(2)\\times U(1)$, \\cite{Farhi:2005rz, Graham:2006vy, Graham:2006xs}. Oscillons first made their appearance in the literature in the 1970s \\cite{Bogolyubsky:1976yu}. They were subsequently rediscovered in the 1990s \\cite{Gleiser:1993pt}. Oscillons are not exact solutions and (very slowly) radiate their energy away. The amplitude of the outgoing radiation (in the small amplitude expansion) has been calculated by a number of authors, see for example \\cite{Segur:1987mg, Fodor:2008es, Fodor:2009kf}. Characterization of their lifetimes and related properties using the ``Gaussian\" ansatz for the spatial profile was done in \\cite{Gleiser:2009ys} (also see references therein). The importance of the dimensionality of space for these objects has been discussed in \\cite{Saffin:2006yk,Gleiser:2004an}. Their possible applications in early Universe physics has not gone unnoticed. For example, they could be relevant for axion dynamics near the QCD phase transition \\cite{Kolb:1993hw}. The properties of oscillons in a $1+1$-dimensional expanding universe (in the small amplitude limit) have been discussed in \\cite{Farhi:2007wj, Graham:2006xs}. Their importance during bubble collisions and phase transitions have been discussed in \\cite{Dymnikova:2000dy}. In \\cite{Hindmarsh:2007jb}, interactions of oscillons with each other and with domain walls were studied in 2+1 dimensions. In this paper, we point out what is required of scalar field potentials to support oscillons. We then derive the frequency as well as the spatial profile of the oscillons for a class of models under consideration. We show that the spatial profile can be very different from a Gaussian, an ansatz often made in the literature. In particular, we derive the nonmonotonic relationship between the height and the width of the oscillons, and discuss the importance of this feature for the stability of oscillons (see \\cite{Lee:1991ax} for a somewhat related analysis for Q-balls). To the best of our knowledge, this has not been done previously in the literature in the context of oscillons. We consider the stability of oscillons against small perturbations, mainly with spatial variations comparable to the width of the oscillons. We also comment on a possible, narrow band instability at higher wave numbers. Oscillons could have important applications in cosmology, especially in the early Universe. With this in mind, we discuss the changes in the profile and the loss of energy from these oscillons due to expansion. The properties of oscillons can depend significantly on the number of spatial dimensions. In this paper, for simplicity we always start with $1+1$-dimensional scenarios where analytic treatment is often possible. We then extend our results to the physically more interesting case of $3+1$ dimensions, analytically where possible and numerically otherwise. We extend previous analysis to the interesting ``flat-top\" oscillons because of our new systematic method for capturing the entire range of possible amplitudes, while still using the methods from the small amplitude expansion. Our expansion, which can also be thought of as a single frequency approximation, allows us to use results existing in the literature for time periodic, localized solutions. Although interesting as classical solutions, a quantum treatment can lead to changes in oscillon lifetimes \\cite{Hertzberg:2009}. Another question worth investigating is the stability of oscillons coupled to other fields. These two questions are beyond the scope of this paper. The rest of the paper is organized as follows: In Sec. \\ref{ScalarOsc} we give a brief overview of oscillons in scalar field theories. In Sec. \\ref{model} we introduce a simple model that is used throughout the paper. Section \\ref{profile} deals with the derivation of the shape and frequency of the oscillons in the absence of expansion. Section \\ref{stability} focuses on linear stability of oscillons. Section \\ref{expansion} discusses the effects of expansion. Our conclusions and future directions are presented in Sec. \\ref{con}. ", "conclusions": "\\label{con} In this paper, we pointed out what is required for scalar field potentials to support oscillons. We found that the spatial profiles can be very different from a Gaussian, an ansatz often made in the literature. In particular, we derived the nonmonotonic relationship between the height and the width of the oscillons, and discussed the importance of this feature for the stability of oscillons. We showed that the flat-topped oscillons are more stable in three dimensions to long wavelength perturbations as compared to their usual Gaussian counterparts. To the best of our knowledge, this had not been done previously in the literature in the context of oscillons. Oscillons could have important applications in cosmology, especially in the early Universe. With this in mind, we discussed the changes in the profile, loss of energy from these oscillons due to expansion, and estimated their lifetimes. We provided analytic results for the $1+1$-dimensional scenario, and extended analytically where possible to $3+1$ dimensions and numerically otherwise. A number of questions related to this work require further investigation. Our expressions for lifetime and arguments for stability, especially in $3+1$ dimensions in the flat-top regimes should be checked with a detailed numerical investigation. The question of the possible small wavelength, narrow band instability needs to be resolved rigorously. Recently, \\cite{Fodor:2009kg} discussed oscillons in the presence of gravity (an oscillaton). It would be interesting to revisit this problem in the context of our large energy, flat-top oscillons. A study of oscillons emerging from (p)reheating-\\cite{Kofman:1997yn,Shtanov:1994ce} like initial conditions in the early Universe is currently in progress." }, "1002/1002.4872_arXiv.txt": { "abstract": "The Sunyaev-Zel'dovich (SZ) effect has a distinct spectral signature that allows its separation from fluctuations in the cosmic microwave background (CMB) and foregrounds. Using CMB anisotropies measured in Wilkinson Microwave Anisotropy Probe's five-year maps, we constrain the SZ fluctuations at large, degree angular scales corresponding to multipoles in the range from 10 to 400. We provide upper bounds on SZ fluctuations at multipoles greater than 50, and find evidence for a hemispherically asymmetric signal at ten degrees angular scales. The amplitude of the detected signal cannot be easily explained with the allowed number density and temperature of electrons in the Galactic halo. We have failed to explain the excess signal as a residual from known Galactic foregrounds or instrumental uncertainties such as $1/f$-noise. ", "introduction": "The Sunyaev-Zel'dovich (SZ) effect~\\cite{Sunyaev:1972eq} is the inverse Compton scattering of the cosmic microwave background (CMB) by electrons throughout the universe. The SZ effect can be partitioned into two components: the kinetic effect (also known as the Ostriker-Vishniac effect~\\cite{Ostriker86}) due to bulk motion of electrons with respect to the rest frame of the microwave background, and the thermal effect due to energy transfer from hot electrons in massive galaxy clusters~\\cite{Komatsu:1999ev, Cooray:2000ge, Molnar:2000de, Springel01, Seljak:2001rc, Sadeh:2004jk}. The SZ thermal signal is expected to be important in future constraints on the underlying cosmology of the universe, by revealing the properties of clusters in cluster-counting experiments (e.g. see Carlstrom et al. 2002~\\cite{Carlstrom:2002na}) and via the angular power spectrum~\\cite{Komatsu:2002wc}. At degree angular scales, there are no detections of the SZ fluctuations as primordial anisotropies of the CMB dominate the SZ signal by at least 3 orders of magnitude. Fortunately, a separation of the SZ anisotropies from those of the CMB is possible due to the fact that the SZ effect has a distinct frequency spectrum, as the inverse-Compton scattering on average increases the net energy of the CMB photons and move photons from the low frequency Rayleigh-Jeans (RJ) tail to higher frequencies~\\cite{Sunyaev:1972eq}. In multi-frequency CMB experiments spanning a wide range of frequency coverage across the SZ null frequency at 217 GHz, one can separate the CMB down to sub-percent level required to study the SZ anisotropy power spectrum~\\cite{Cooray:2000xh}. This approach was first applied to constrain the sub-degree SZ fluctuations in BOOMERanG 2003 data~\\cite{Veneziani:2009es}. Here, we analyze the WMAP five-year CMB data to constrain SZ anisotropies at degree angular scales corresponding to multipoles $\\ell=10-400$ on the sky. We ignore $\\ell < 10$ as a measure to avoid biases from previously reported anomalies, the divergence of 1/$f$-noise, and large uncertainties associated with simulating the CMB sky properly to match WMAP data at these multipoles~\\cite{ben92,hin96,copi04,land05,schwarz04,oliveira04,eriksen04,bennett10}. In Section 2 we detail our approach to extract the SZ signal from the multifrequency WMAP5 data set; in Section 3 we present our results; and in Section 4 we provide a discussion of these results. ", "conclusions": "\\label{sec:discussion} If the excess signal we have detected at ten degrees angular scales is indeed SZ, it is clear that the signal does not originate from galaxy clusters at high redshift. Such a signal would lead to a power spectrum that increases with multipole, rendering an SZ signal in the smaller scale bins. A large SZ signal associated with the Extragalactic sky is also ruled out by existing arcminute-scale CMB experiments. We do not see the signal coming from unresolved point sources or Zodiacal light emission as they do not amount to more than a few $\\mu$K in strength~\\cite{Diego:2009wi}. One possibility is that the signal is associated with sub-keV gas in the Galactic halo, explaining why the signal is primarily at small multipoles. We can try to address if such a signal is possible based on requirements on the number density of electrons and the existing spectral distortion limits from FIRAS. Since the SZ effect is $-2y$ at RJ frequencies, where $y$ is the Compton $y$-parameter, we can convert the amplitude of the fluctuation power, for a Galactic-like signal with power-law $C_{\\ell} \\simeq A \\ell^{-3}$~\\cite{Lagache2007} and assuming a dispersion in $y$ equal to its average value, to obtain $y = (5.1 \\pm 0.8) \\times 10^{-6}$ (95\\% c.l.). This is consistent with the 95\\%-level FIRAS constraint of $|y| \\leq 1.5 \\times 10^{-5}$ from the COBE satellite~\\cite{Fixsen:1996nj}. The electron column density required to render this $y$-parameter in our Milky Way halo can be obtained from $N_e \\simeq {{y m_ec^2} \\over {\\sigma_T k_{\\rm B}T_e}}$, where $m_e$ is the electron mass, $T_e$ is the electron temperature, with other constants having the usual definition. For a 0.3 keV electron gas temperature, we find $N_e = {1.3 \\pm 0.2} \\times 10^{22}~{\\rm cm}^{-2} \\left({{0.3~{\\rm keV}} \\over {k_{\\rm B}T_e}}\\right)$ (95\\% c.l.). This electron density is merely approximate due to the assumptions made in deriving $y$, and in presuming all of the measured signal is actually rooted in the SZ. The column density of the free electrons is consistent with that suggested in Peiris \\& Smith (2010)~\\cite{Peiris:2010jd} to resolve the CMB isotropy anomalies via the kinetic SZ effect. In reality, both the thermal and kinetic SZ effects may contribute to large-scale anomalies of the CMB. However, the required column density is in strong tension with pulsar dispersion measurements (Taylor \\& Cordes 1993~\\cite{Taylor:1993my}), OVII absorption studies~\\cite{Fang:2010yk}, and the soft X-ray background~\\cite{Yoshino:2009kv}, which place $N_e \\sim {10^{20}-10^{21}} \\rm{cm}^{-2}$. Given the WMAP subtracted foregrounds would need to have been underestimated by a factor of two to explain our low-$\\ell$ signal, this discrepancy suggests that the excess signal may originate from further instrumental effects or an unknown foreground. Rather than to simply explain away the signal, as we did not find either a simple statistical argument or systematic effect, we suggest that further studies are warranted." }, "1002/1002.1444_arXiv.txt": { "abstract": "Stereo data collected by the HiRes experiment over a six year period are examined for large-scale anisotropy related to the inhomogeneous distribution of matter in the nearby Universe. We consider the generic case of small cosmic-ray deflections and a large number of sources tracing the matter distribution. In this matter tracer model the expected cosmic ray flux depends essentially on a single free parameter, the typical deflection angle $\\theta_s$. We find that the HiRes data with threshold energies of 40 EeV and 57 EeV are incompatible with the matter tracer model at a 95\\% confidence level unless $\\theta_s>10^\\circ$ and are compatible with an isotropic flux. The data set above 10 EeV is compatible with both the matter tracer model and an isotropic flux. ", "introduction": "The observation of the cutoff in the spectrum of Ultra-High Energy Cosmic Rays (UHECRs) \\citep{Abbasi:2007sv,Abraham:2008ru} as predicted by \\citet{Greisen:1966jv,Zatsepin:1966jv} provides compelling evidence for the shortening of the UHECR propagation length at high energies. The highest energy events then must have come from relatively close sources (within $250$ Mpc). At these length scales the matter in the Universe is distributed inhomogeneously, being organized into clusters and superclusters. One should, therefore, expect the flux of highest-energy cosmic rays to be anisotropic. In astrophysical scenarios, it is natural to assume that the number of sources within $250$~Mpc is large, and that these sources trace the distribution of matter. Under these assumptions, the anisotropy at Earth depends only on the nature and size of UHECR deflections. Measurement of the anisotropy, therefore, gives direct experimental access to parameters that determine the deflections, notably to the UHECR charge composition and cosmic magnetic fields. Several investigations of anisotropy in arrival directions of UHECRs have been previously undertaken. At small angular scales, correlations with different classes of putative sources were claimed (e.g. \\citealt{Gorbunov:2004bs,Abbasi:2005qy,Cronin:2007zz,Abraham:2007si}). At larger angular scales and energies below 10 EeV possible anisotropy towards the Galactic center was reported in \\citet{Hayashida:1998qb,Hayashida:1999ab,Bellido:2000tr}, but not supported by more recent studies \\citep{Santos:2007na}. At higher energies, \\citet{Stanev:1995my} found evidence against an isotropic flux above 40 EeV through correlations with the supergalactic plane, but this was not confirmed by other authors \\citep{Hayashida:1996bc,Kewley:1996zt, Bird:1998nu}. Finally, using the Pierre Auger Observatory (PAO) data, \\citet{Kashti:2008bw} have found correlations between UHECR arrival directions and the large-scale structure of the Universe which are incompatible with an isotropic flux (see, however, \\citealt{Koers:2008ba}). In this paper, we analyze the data accumulated by the HiRes experiment for anisotropy associated with the large-scale structure of the Universe. The HiRes experiment has been described previously \\citep{HiResStatus1999,Boyer:2002zz,Hanlon:2008}. It studied ultrahigh energy cosmic rays from $10^{17.2}$ eV to $10^{20.2}$ eV using the fluorescence technique. HiRes operated two fluorescence detectors located atop desert mountains in west-central Utah. The data set used in this study consists of events observed by both detectors, analyzed jointly in what is commonly called ``stereo mode''. In this mode the angular resolution in cosmic rays' pointing directions is about $0.8$ degrees, and the energy resolution is about 10\\%. The HiRes experiment operated in stereo mode between December, 1999, and April, 2006. At the highest energies HiRes has the largest data set in the Northern hemisphere. Large number of events, good angular resolution (better than the bending angles expected from Galactic and extragalactic magnetic fields) and the wide energy range covered make the HiRes data particularly suitable for anisotropy studies. The exact data set used in this study was described previously in \\citet{Abbasi:2008md}. We consider here a generic model that assumes many sources within $250$~Mpc tracing the distribution of matter, which we refer to as the ``matter tracer'' model. We also assume that deflections of UHECR do not exceed the angular size of the nearby structures, that is 10-20$^\\circ$. In this regime, both regular and random deflections in magnetic fields can be modeled with a one-parameter distribution, for which we take a Gaussian distribution centered at zero with width $\\theta_{\\rm s}$. This width is treated as a free parameter, whose value we aim to constrain from the data. Constraints on $\\theta_{\\rm s}$ may then be used to obtain information on the strength of Galactic and extragalactic magnetic fields. In keeping with our assumption of small deflections, we assume a proton composition in this study, which is consistent with the $X_{\\rm max}$ analysis based on the same dataset (for confirmation see \\citealt{Abbasi:2009nf}). The HiRes data is compared to model predictions using the ``flux sampling'' test put forward by \\citet{Koers:2008ba}. This test has good discrimination power at small statistics and is insensitive to the details of deflections. The comparison is performed at three different threshold energies that have been used in previous studies: 10 EeV, 40 EeV, and 57 EeV \\citep{Hayashida:1996bc,Abbasi:2005qy,Cronin:2007zz}. An {\\em a priori} significance of 5\\%, corresponding to a confidence level (CL) of 95\\%, is chosen for this work. The paper is organized as follows. In section~\\ref{section:modeling} we discuss the modeling of UHECR arrival directions. Section~\\ref{section:data} concerns the HiRes data used in the analysis, while section~\\ref{section:fluxsampling} describes the flux sampling method. We present our results in section~\\ref{sec:results} and conclude in section~\\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} To summarize, we have confronted the stereo data collected by the HiRes experiment with predictions of the matter tracer model, a generic model of cosmic ray origin and propagation. The model assumes a large number of cosmic ray sources within $250$~Mpc whose distribution traces that of matter, and relatively small deflections characterized by a single parameter, the typical deflection angle $\\theta_s$. We have found that the HiRes data with energy thresholds $E=40$~EeV and $E=57$~EeV are incompatible with the matter tracer model for $\\theta_s<10^\\circ$ at 95\\%~CL. With an energy threshold $E=10$~EeV the HiRes data are compatible with the matter tracer model. At all three energy thresholds, the data are compatible with an isotropic flux at 95\\%~CL. In the present analysis we have treated the deflections as random and Gaussian, which is only appropriate for small deflection angles and limited number of events. The actual deflections are expected to contain a coherent component due to the Galactic magnetic field. With the accumulation of UHECR events by PAO in the Southern hemisphere and by Telescope Array in the Northern hemisphere, our analysis will become sensitive to the nature of deflections and thus, with proper modifications of the statistical procedure, may give direct access to the parameters of cosmic magnetic fields." }, "1002/1002.4866_arXiv.txt": { "abstract": "The addition of Wide Field Camera 3 (WFC3) on the {\\em Hubble Space Telescope} ({\\em HST}) has led to a dramatic increase in our ability to study the $z>6$ Universe. The improvement in the near-infrared (NIR) sensitivity of WFC3 over previous instruments has enabled us to reach apparent magnitudes approaching $29$ (AB). This allows us to probe the rest-frame ultraviolet (UV) continuum, redshifted into the NIR at $z>6$. Taking advantage of the large optical depths of the intergalactic medium at this redshift, resulting in the Lyman-$\\alpha$ break, we use a combination of WFC3 imaging and pre-existing Advanced Camera for Surveys (ACS) imaging to search for $z\\approx 7$ galaxies over 4 fields in and around Great Observatories Origins Survey (GOODS) South. Our analysis reveals 29 new $z\\approx 7$ star forming galaxy candidates in addition to 15 pre-existing candidates already discovered in these fields. The improved statistics from our doubling of the robust sample of $z$-drop candidates confirms the previously observed evolution of the bright end of the luminosity function. ", "introduction": "There has been great progress in recent years in discovering star-forming galaxies at high redshift ($z>5$), either through the Lyman-$\\alpha$ emission line or their rest-frame UV continuum colours, which are strongly affected by absorption by intervening hydrogen (the Gunn-Peterson effect at $z>6.3$ provides almost complete absorption shortward of Lyman-$\\alpha$, e.g. Gunn \\& Peterson 1965, Fan et al. 2001). This Lyman-break technique was initially used to identify $z\\approx 3$ galaxies from broad-band images (Steidel et al.\\ 1996), and improvements in sensitivity of near-infrared imaging have enabled this to be pushed to higher redshifts (e.g. at $z\\approx 6$ using the Advanced Camera for Surveys on {\\em HST}, Stanway, Bunker \\& McMahon 2003). Recently, the new {\\em Wide Field Camera 3} (WFC3) camera has been installed on {\\em HST}, which has a near-infrared channel with significantly larger field and better sensitivity than the previous-generation {\\em Near Infrared Camera and Multi-Object Spectrometer} (NICMOS) instrument. A number of multi-waveband imaging campaigns are underway, typically using broad-band filters at 1.0, 1.25 and 1.6$\\mu$m and targetting fields with existing deep optical data. This is an ideal strategy to look for optical ``drop-outs'', seen only at the longer WFC3 wavelengths, which are candidate $z>7$ galaxies. A number of papers (e.g. Oesch et al. 2010, Bouwens et al. 2010, Bunker et al. 2010, McLure et al. 2010, Yan et al. 2010 and Finkelstein et al. 2010, Wilkins et al. 2010) have appeared in the past year presenting first results from such imaging. In this paper we analyse all the data taken so far in two large programs targetting fields in the vicinity of the Great Observatories Orgins Deep Survey-South (GOODS-South), which in one case includes the {\\em Hubble} Ultra Deep Field, and present a new sample of $z$-band drop-outs which more than doubles the existing number of robust $z>7$ candidates. By building such large samples ($>40$ objects spanning 3 magnitudes) we can ultimately study the rest-UV luminosity function, and hence estimate the integrated star formation rate within 800\\,Myr of the Big Bang. Understanding the global star formation history at these redshifts is crucial in answering the question whether the UV photons produced by the short-lived massive OB stars were sufficient to reionize the Universe. This paper is organised as follows: in Section 2 we discuss the observations, our data reduction process and our photometry measurement and catalogue production technique. In Section 3 our candidate selection procedure is described and we present a list of $z$-drop candidates together with a comparison with other studies. In Section 4 we investigate the properties of these objects, including their redshift distribution, surface density, and luminosity function. Further, in this section we investigate the implications of the luminosity function for the star formation rate density and production of ionising photons. In Section 5 we present our conclusions. Throughout, we adopt the standard \u2018concordance\u2019 cosmology of $\\Omega_{M}= 0.3$, $\\Omega_{\\Lambda}= 0.7$ and use $H_{0}=70$ kms$^{-1}$ Mpc$^{-1}$. All magnitudes are on the AB system (Oke \\& Gunn 1983). ", "conclusions": "In this work we have presented a catalogue of candidate $z$-drop $z\\sim 7$ galaxies from 4 separate fields (HUDF, P34, P12, ERS) covering a total area of $\\sim 50.0$ arcmin$^{2}$. The deepest field, the HUDF, extends to $J_{125w}=28.36$ ($7\\sigma$) and shallowest, the ERS, $J_{125w}=27.03$ ($7\\sigma$). In total we present 44 candidates (HUDF:11, P34:15, P12:7, ERS:11) covering an apparent magnitude range of $J_{125w}=25.9-28.36$. Simulations reveal these galaxies likely cover a redshift distribution $z=6.4-7.7$. Of these 44 candidates, 15 are recorded in a range of previous studies. While our selection criteria guards against any intrinsic interloper population it is still susceptible to contamination through photometric scatter and transient phenomena (supernovae and significant apparent motion objects). Simulations testing the effect of photometric scatter suggest a contamination rate of $\\sim 5\\%$ ($\\sim 2.5$ objects). Further, we note the presence of a possible supernovae in the HUDF field and record the presence of an object, which due to significant proper motion, mimics the characteristics of a high-redshift galaxy. Using this new data set, we determine the stepwise rest-frame ultraviolet luminosity function for each field ($-21.0\\alpha>-1.9$, motivated by observations at lower redshift, the remaining luminosity function parameters can be determined. From these the ultraviolet luminosity density, and thus both the star formation rate and ionising photon luminosity densities can be estimated. The ionising photon luminosity density inferred from the observed luminosity function extrapolated to $M_{1600}=-8.0$ assuming either a steep steep faint-end $\\alpha<-1.86$ or a shallower slope with a lower metallicity and/or different IMF is sufficient to maintain the ionisation of the Universe for recent estimates of the hydrogen clumping factor and ionising photon escape fraction. \\subsection*{Acknowledgements} We would like to thank the referee for their detailed and helpful suggestions. SMW acknowledges the support of STFC and SL and JC acknowledge the support of ELIXIR. Based on observations made with the NASA/ESA Hubble Space Telescope, obtained from the Data Archive at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. These observations are associated with programme \\#GO/DD-11359 and \\#GO-11563. We are grateful to the WFC\\,3 Science Oversight Committee for making their Early Release Science observations public." }, "1002/1002.1996_arXiv.txt": { "abstract": "We present an occultation of the newly discovered hot Jupiter system WASP-19, observed with the HAWK-I instrument on the VLT, in order to measure thermal emission from the planet's dayside at $\\sim$2$\\umu$m. The light curve was analysed using a Markov-Chain Monte-Carlo method to find the eclipse depth and the central transit time. The transit depth was found to be 0.366$\\pm$0.072\\%, corresponding to a brightness temperature of 2540$\\pm$180\\,K. This is significantly higher than the calculated (zero-albedo) equilibrium temperature, and indicates that the planet shows poor redistribution of heat to the night side, consistent with models of highly irradiated planets. Further observations are needed to confirm the existence of a temperature inversion, and possibly molecular emission lines. The central eclipse time was found to be consistent with a circular orbit. ", "introduction": " ", "conclusions": "We have detected an occultation of the extrasolar planet system WASP-19 in the NB2090 filter centred at 2.095\\,$\\umu$m using the HAWK-I instrument. The eclipse depth was measured to be 0.366$\\pm$0.072\\%, with the transit centre occurring at phase 0.50114$\\pm$0.00241, consistent with a circular orbit to well within 1$\\sigma$. \\citet{anderson_2010}, however, report evidence for a non-zero eccentricity at the 2.6$\\sigma$ level. The methods used by \\citet{anderson_2010} to measure the phase of the eclipse, and to correct for systematics, particularly the trends in the light curve which can considerably effect transit times \\citep[e.g.][]{gibson_2009}, differs from that adopted in our analysis, and this difference is the most likely source of the apparent discrepancy. The brightness temperature may be calculated from the depth, taking the values and uncertainties for the stellar temperature and planet-to-star radius ratio from \\citet{hebb_2010}. This results in a brightness temperature of 2540$\\pm$180\\,K, considerably larger than the (zero-albedo) equilibrium temperature given in \\citet{hebb_2010} of 2009$\\pm$26\\,K. \\citet{anderson_2010} report an eclipse depth of $0.259^{+0.046}_{-0.044}\\%$ in the H-band, corresponding to a brightness temperature of 2560$\\pm$130\\,K. This is again considerably higher than the equilibrium temperature, and also the equilibrium temperature given no redistribution of heat to the nightside, which they calculate as $\\sim2400$\\,K. Our 2.095\\,$\\umu$m measurement is consistent with this temperature. This indicates that there is poor redistribution of heat to the night side of the planet, consistent with the pM class of planets in the classification of \\citet{fortney_2008}. They also conclude that planets hot enough to have significant TiO and VO absorption in their upper atmospheres would show temperature inversions, and may appear anomalously bright in the infrared due to molecular emission bands, which may help explain the temperature excesses here. Other ground based K-band secondary eclipse measurements of very hot Jupiters reached similar conclusions \\citep{demooij_snellen,gillon_2009,alonso_2010}. Further observations at infrared wavelengths are required to confirm temperature inversions and possibly measure molecular emission for WASP-19, which should prove an interesting target for future studies." }, "1002/1002.3488_arXiv.txt": { "abstract": "Name{ABSTRACT}% \\def\\abstract{% \\endmode \\bigskip\\bigskip \\centerline{\\AbstractName}% \\medskip \\bgroup \\let\\endmode=\\endabstract \\narrower\\narrower \\singlespaced \\everyabstract}% \\def\\everyabstract{}% \\def\\endabstract{\\smallskip\\egroup} \\def\\pacs#1{\\medskip\\centerline{PACS numbers: #1}\\smallskip} \\def\\submit#1{\\bigskip\\centerline{Submitted to {\\sl #1}}} \\def\\submitted#1{\\submit{#1}}% \\def\\toappear#1{\\bigskip\\raggedcenter To appear in {\\sl #1} \\endraggedcenter} \\def\\disclaimer#1{\\footnote{}\\bgroup\\tenrm\\singlespaced This manuscript has been authored under contract number #1 \\@disclaimer\\par} \\def\\disclaimers#1{\\footnote{}\\bgroup\\tenrm\\singlespaced This manuscript has been authored under contract numbers #1 \\@disclaimer\\par} \\def\\@disclaimer{% with the U.S. Department of Energy. 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\\let\\@FigInsert=#1\\relax \\def\\@arg{#2}\\ifx\\@arg\\empty\\def\\@ID{}% \\else\\ifRomanTables \\global\\advance\\@count by\\@ne \\edef\\@ID{\\uppercase\\expandafter {\\romannumeral\\the\\@count}}% \\tag{\\@prefix#2}{\\@ID}% \\else \\LabelParse #2;;\\endlist\\fi \\fi \\ifTabsLast \\emsg{\\CaptionName\\space\\@ID. {#2} [storing in \\jobname.tb]}% \\@tbwrite{\\@comment> \\CaptionName\\space\\@ID.\\space{#2}}% \\@tbwrite{\\string\\@FigureItem{\\CaptionName}{\\@ID}{\\NX#1}}% \\seeCR\\let\\@next=\\@copytab \\else \\emsg{\\CaptionName\\ \\@ID.\\ {#2}}% \\let\\endtable=\\@endfigure \\setbox\\@capbox\\vbox to 0pt{}% \\def\\@whereCap{N}% \\let\\@next=\\@findcap \\ifx\\@FigInsert\\midinsert\\goodbreak\\fi \\@FigInsert \\fi \\@next} \\def\\endtable{\\emsg{> \\string\\endtable before \\string\\table!}} \\def\\@copytab#1#2\\endtable{\\endgroup \\ifx#1\\par\\@tbNXwrite{#2\\@endfigure}\\else\\@tbNXwrite{#1#2\\@endfigure}\\fi} \\def\\@TBinit{\\@FileInit\\tbout=\\jobname.tb[Tables]\\gdef\\@TBinit{\\relax}} \\def\\@tbwrite#1{\\@TBinit\\immediate\\write\\tbout{#1}} \\long\\def\\@tbNXwrite#1{\\@TBinit\\unexpandedwrite\\tbout{#1}} \\def\\PrintTables{\\ifTabsLast\\@PrintTables\\fi} \\def\\@PrintTables{% \\@tbwrite{\\@comment>>> EOF \\jobname.tb <<<}% \\immediate\\closeout\\tbout \\begingroup \\TabsLastfalse \\catcode`@=11 \\offparens \\emsg{[Getting tables from file \\jobname.tb]}% \\Input\\jobname.tb \\relax \\endgroup}% \\newbox\\@capbox \\newcount\\@caplines \\def\\CaptionName{}% \\def\\@ID{}% \\def\\captionspacing{\\normalbaselines}% \\def\\@inCaption{F}% \\long\\def\\caption#1{% \\def\\lab@l{\\@ID}% \\global\\setbox\\@capbox=\\vbox\\bgroup \\def\\@inCaption{T}% \\captionspacing\\seeCR \\dimen@=20\\parindent \\ifdim\\colwidth>\\dimen@\\narrower\\narrower\\fi \\noindent{\\bf \\linkname{\\@TagName}{\\CaptionName~\\@ID}:\\space}% #1\\relax \\vskip 0pt \\global\\@caplines=\\prevgraf \\egroup \\ifnum\\@caplines=\\@ne \\global\\setbox\\@capbox=\\vbox{\\noindent\\seeCR \\hfil{\\bf \\linkname{\\@TagName}{\\CaptionName~\\@ID}:\\space}% #1\\hfil}\\fi \\if N\\@whereCap\\def\\@whereCap{B}\\fi \\if T\\@whereCap \\centerline{\\box\\@capbox}% \\vskip\\baselineskip\\medskip \\fi}% \\def\\Caption{\\begingroup\\seeCR\\@Caption}% \\long\\def\\@Caption#1\\endCaption{\\endgroup \\ifCaptionList \\incaplist{#1}\\fi \\caption{#1}}% \\def\\endCaption{\\emsg{> \\string\\endCaption\\ called before \\string\\Caption.}} \\long\\def\\@findcap#1{% \\ifx #1\\Caption \\def\\@whereCap{T}\\fi \\ifx #1\\caption \\def\\@whereCap{T}\\fi #1}% \\def\\@whereCap{N}% \\def\\ListCaptions{\\@ListCaps\\caplist=\\jobname.cap[List of Captions] {\\let\\FIGLitem=\\CAPLitem}} \\def\\ListFigureCaptions{% \\@ListCaps\\figlist=\\jobname.fgl[List of Figure Captions] {\\let\\FIGLitem=\\CAPLitem}} \\def\\ListTableCaptions{% \\@ListCaps\\tablelist=\\jobname.tbl[List of Table Captions] {\\let\\FIGLitem=\\CAPLitem}} \\def\\CAPLitem#1#2#3\\@endFIGLitem#4{% \\bigskip \\begingroup \\raggedright\\tolerance=1700 \\hangindent=1.41\\parindent\\hangafter=1 \\noindent #1\\ #2 #3 \\hskip 0pt plus 10pt \\vskip 0pt \\endgroup}% \\def\\infiglist{\\begingroup\\seeCR \\@infiglist\\figlist} \\def\\intablelist{\\begingroup\\seeCR \\@infiglist\\tablelist} \\def\\incaplist{\\begingroup\\seeCR \\@infiglist\\caplist} \\def\\FigListWrite#1#2{% \\ifx#1\\figlist\\relax \\FigListInit\\fi \\ifx#1\\tablelist\\relax \\TabListInit\\fi \\ifx#1\\caplist\\relax \\CapListInit\\fi \\edef\\@line@{{#2}}% \\write#1\\@line@}% \\def\\FigListInit{\\@FileInit\\figlist=\\jobname.fgl[List of Figures]% \\gdef\\FigListInit{\\relax}} \\def\\TabListInit{\\@FileInit\\tablelist=\\jobname.tbl[List of Tables]% \\gdef\\TabListInit{\\relax}} \\def\\CapListInit{\\@FileInit\\caplist=\\jobname.cap[List of Captions]% \\gdef\\CapListInit{\\relax}} \\def\\FigListWriteNX#1#2{\\writeNX#1{#2}} \\def\\@infiglist#1#2{% \\FigListWrite#1{\\@comment > \\CaptionName\\space\\@ID:}% \\FigListWrite#1{\\string\\FIGLitem{\\CaptionName} {\\@ID.\\space}}% \\@copycap#1#2\\endlist \\FigListWrite#1{{\\NX\\folio}}% \\endgroup}% \\def\\@copycap#1#2#3\\endlist{% \\ifx#2\\space\\writeNX#1{#3\\@endFIGLitem}% \\else\\writeNX#1{#2#3\\@endFIGLitem}\\fi} \\def\\ListFigures{\\@ListCaps\\figlist=\\jobname.fgl[List of Figures]{}} \\def\\ListTables{\\@ListCaps\\tablelist=\\jobname.tbl[List of Tables]{}} \\def\\@ListCaps#1=#2[#3]#4{% \\immediate\\closeout#1 \\openin#1=#2 \\relax \\ifeof#1\\closein#1 \\else\\closein#1\\emsg{[Getting #3]}% \\begingroup \\showsectIDtrue \\ATunlock\\quoteoff\\offparens #4 \\input #2 \\relax \\endgroup \\fi} \\long\\def\\FIGLitem#1#2#3\\@endFIGLitem#4{% \\medskip \\begingroup \\raggedright\\tolerance=1700 \\ifx\\TOCmargin\\undefined\\skip0=\\parindent \\else\\skip0=\\TOCmargin\\fi \\advance\\rightskip by \\skip0 \\parfillskip=-\\skip0 \\hangindent=1.41\\parindent\\hangafter=1 \\noindent \\ifshowsectID #1\\ \\fi #2 #3 \\hskip 0pt plus 10pt \\leaddots \\hbox to 2em{\\hss\\linkto{page.#4}{#4}}% \\vskip 0pt \\endgroup} \\def\\Fig#1{\\linkto{Fg.#1}{Fig.~\\use{Fg.#1}}} \\def\\Figs#1{\\linkto{Fg.#1}{Figs.~\\use{Fg.#1}}} \\def\\Fg#1{\\linkto{Fg.#1}{\\use{Fg.#1}}} \\def\\Tab#1{\\linkto{Tb.#1}{Table~\\use{Tb.#1}}} \\def\\Tbl#1{\\linkto{Tb.#1}{Table~\\use{Tb.#1}}} \\def\\Tb#1{\\linkto{Tb.#1}{\\use{Tb.#1}}} \\autoload\\Tablebody{Tablebod.txs}\\autoload\\Tablebodyleft{Tablebod.txs} \\autoload\\tablebody{Tablebod.txs} \\autoload\\epsffile{epsf.tex} \\autoload\\epsfbox{epsf.tex} \\autoload\\epsfxsize{epsf.tex} \\autoload\\epsfysize{epsf.tex} \\autoload\\epsfverbosetrue{epsf.tex}\\autoload\\epsfverbosefalse{epsf.tex} \\obsolete\\topFigure\\figure \\obsolete\\midFigure\\midfigure \\obsolete\\fullFigure\\fullfigure \\obsolete\\TOPFIGURE\\figure \\obsolete\\MIDFIGURE\\midfigure \\obsolete\\FULLFIGURE\\fullfigure \\obsolete\\endFigure\\endfigure \\obsolete\\ENDFIGURE\\endfigure \\obsolete\\topTable\\toptable \\obsolete\\midTable\\midtable \\obsolete\\fullTable\\fulltable \\obsolete\\TOPTABLE\\toptable \\obsolete\\MIDTABLE\\midtable \\obsolete\\FULLTABLE\\fulltable \\obsolete\\endTable\\endtable \\obsolete\\ENDTABLE\\endtable \\def\\FIG{\\@obsolete\\FIG\\Fig\\Fig}% \\def\\TBL{\\@obsolete\\TBL\\Tbl\\Tbl}% \\catcode`@=11 \\catcode`\\|=12 \\catcode`\\&=4 \\newcount\\ncols \\ncols=\\z@ \\newcount\\nrows \\nrows=\\z@ \\newcount\\curcol \\curcol=\\z@ \\let\\currow=\\nrows \\newdimen\\thinsize \\thinsize=0.6pt \\newdimen\\thicksize \\thicksize=1.5pt \\newdimen\\tablewidth \\tablewidth=-\\maxdimen \\newdimen\\parasize \\parasize=4in \\newif\\iftableinfo \\tableinfotrue \\newif\\ifcentertables \\centertablestrue \\def\\centeredtables{\\centertablestrue}% \\def\\noncenteredtables{\\centertablesfalse}% \\def\\nocenteredtables{\\centertablesfalse}% \\let\\plaincr=\\cr \\let\\plainspan=\\span \\let\\plaintab=& \\def\\ampersand{\\char`\\&}% \\let\\lparen=( \\let\\NX=\\noexpand \\def\\ruledtable{\\relax \\@BeginRuledTable \\@RuledTable}% \\def\\@BeginRuledTable{% \\ncols=0\\nrows=0 \\begingroup \\offinterlineskip \\def~{\\phantom{0}}% \\def\\span{\\plainspan\\omit\\relax\\colcount\\plainspan}% \\let\\cr=\\crrule \\let\\CR=\\crthick \\let\\nr=\\crnorule \\let\\|=\\Vb \\def\\hfill{\\hskip0pt plus1fill\\hbox{}}% \\ifx\\tablestrut\\undefined\\relax \\else\\let\\tstrut=\\tablestrut\\fi \\catcode`\\|=13 \\catcode`\\&=13\\relax \\TableActive \\curcol=1 \\ifdim\\tablewidth>-\\maxdimen\\relax \\edef\\@Halign{\\NX\\halign to \\NX\\tablewidth\\NX\\bgroup\\TablePreamble}% \\tabskip=0pt plus 1fil \\else \\edef\\@Halign{\\NX\\halign\\NX\\bgroup\\TablePreamble}% \\tabskip=0pt \\fi \\ifcentertables \\ifhmode\\vskip 0pt\\fi \\line\\bgroup\\hss \\else\\hbox\\bgroup \\fi}% \\long\\def\\@RuledTable#1\\endruledtable{% \\vrule width\\thicksize \\vbox{\\@Halign \\thickrule #1\\killspace \\tstrut \\linecount \\plaincr\\thickrule \\egroup}% \\vrule width\\thicksize \\ifcentertables\\hss\\fi\\egroup \\endgroup \\global\\tablewidth=-\\maxdimen \\iftableinfo \\immediate\\write16{[Nrows=\\the\\nrows, Ncols=\\the\\ncols]}% \\fi}% \\def\\TablePreamble{% \\TableItem{####}% \\plaintab\\plaintab \\TableItem{####}% \\plaincr}% \\def\\@TableItem#1{% \\hfil\\tablespace #1\\killspace \\tablespace\\hfil }% \\def\\@tableright#1{% \\hfil\\tablespace\\relax #1\\killspace \\tablespace\\relax}% \\def\\@tableleft#1{% \\tablespace\\relax #1\\killspace \\tablespace\\hfil}% \\let\\TableItem=\\@TableItem \\def\\RightJustifyTables{\\let\\TableItem=\\@tableright}% \\def\\LeftJustifyTables{\\let\\TableItem=\\@tableleft}% \\def\\NoJustifyTables{\\let\\TableItem=\\@TableItem}% \\def\\LooseTables{\\let\\tablespace=\\quad}% \\def\\TightTables{\\let\\tablespace=\\space}% \\LooseTables \\def\\TrailingSpaces{\\let\\killspace=\\relax}% \\def\\NoTrailingSpaces{\\let\\killspace=\\unskip}% \\TrailingSpaces \\def\\setRuledStrut{% \\dimen@=\\baselineskip \\advance\\dimen@ by-\\normalbaselineskip \\ifdim\\dimen@<.5ex \\dimen@=.5ex\\fi \\setbox0=\\hbox{\\lparen}% \\dimen1=\\dimen@ \\advance\\dimen1 by \\ht0% \\dimen2=\\dimen@ \\advance\\dimen2 by \\dp0% \\def\\tstrut{\\vrule height\\dimen1 depth\\dimen2 width\\z@}% }% \\def\\tstrut{\\vrule height 3.1ex depth 1.2ex width 0pt}% \\def\\bigitem#1{% \\dimen@=\\baselineskip \\advance\\dimen@ by-\\normalbaselineskip \\ifdim\\dimen@<.5ex \\dimen@=.5ex\\fi \\setbox0=\\hbox{#1}% \\dimen1=\\dimen@ \\advance\\dimen1 by \\ht0 \\dimen2=\\dimen@ \\advance\\dimen2 by \\dp0 \\vrule height\\dimen1 depth\\dimen2 width\\z@ \\copy0}% \\def\\vctr#1{\\hfil\\vbox to 0pt{\\vss\\hbox{#1}\\vss}\\hfil}% \\def\\nextcolumn#1{% \\plaintab\\omit#1\\relax\\colcount \\plaintab}% \\def\\tab{% \\nextcolumn{\\relax}}% \\let\\novb=\\tab \\def\\vb{% \\nextcolumn{\\vrule width\\thinsize}}% \\def\\Vb{% \\nextcolumn{\\vrule width\\thicksize}}% \\def\\dbl{% \\nextcolumn{\\vrule width\\thinsize \\hskip 2\\thinsize \\vrule width\\thinsize}}% {\\catcode`\\|=13 \\let|0 \\catcode`\\&=13 \\let&0 \\gdef\\TableActive{\\let|=\\vb \\let&=\\tab}% }% \\def\\crrule{\\killspace \\tstrut \\linecount \\plaincr\\tablerule }% \\def\\crthick{\\killspace \\tstrut \\linecount \\plaincr\\thickrule }% \\def\\crnorule{\\killspace \\tstrut \\linecount \\plaincr }% \\def\\crpart{\\killspace \\linecount \\plaincr}% \\def\\tablerule{\\noalign{\\hrule height\\thinsize depth 0pt}}% \\def\\thickrule{\\noalign{\\hrule height\\thicksize depth 0pt}}% \\def\\cskip{\\omit\\relax}% \\def\\crule{\\omit\\leaders\\hrule height\\thinsize depth0pt\\hfill}% \\def\\Crule{\\omit\\leaders\\hrule height\\thicksize depth0pt\\hfill}% \\def\\linecount{% \\global\\advance\\nrows by1 \\ifnum\\ncols>0 \\ifnum\\curcol=\\ncols\\relax\\else \\immediate\\write16 {\\NX\\ruledtable warning: Ncols=\\the\\curcol\\space for Nrow=\\the\\nrows}% \\fi\\fi \\global\\ncols=\\curcol \\global\\curcol=1}% \\def\\colcount{\\relax \\global\\advance\\curcol by 1\\relax}% \\long\\def\\para#1{% \\vtop{\\hsize=\\parasize \\normalbaselines \\noindent #1\\relax \\vrule width 0pt depth 1.1ex}% }% \\def\\begintable{\\relax \\@BeginRuledTable \\@begintable}% \\long\\def\\@begintable#1\\endtable{% \\@RuledTable#1\\endruledtable}% \\def\\E#1{\\hbox{$\\times 10^{#1}$}} \\def\\square{\\hbox{{$\\sqcup$}\\llap{$\\sqcap$}}}% \\def\\grad{\\nabla}% \\def\\del{\\partial}% \\def\\frac#1#2{{#1\\over#2}} \\def\\smallfrac#1#2{{\\scriptstyle {#1 \\over #2}}} \\def\\half{\\ifinner {\\scriptstyle {1 \\over 2}}% \\else {\\textstyle {1 \\over 2}}\\fi} \\def\\bra#1{\\langle#1\\vert}% \\def\\ket#1{\\vert#1\\/\\rangle}% \\def\\vev#1{\\langle{#1}\\rangle}% \\def\\simge{% \\mathrel{\\rlap{\\raise 0.511ex \\hbox{$>$}}{\\lower 0.511ex \\hbox{$\\sim$}}}} \\def\\simle{% \\mathrel{\\rlap{\\raise 0.511ex \\hbox{$<$}}{\\lower 0.511ex \\hbox{$\\sim$}}}} \\def\\gtsim{\\simge}% \\def\\ltsim{\\simle}% \\def\\therefore{% \\setbox0=\\hbox{$.\\kern.2em.$}\\dimen0=\\wd0 \\mathrel{\\rlap{\\raise.25ex\\hbox to\\dimen0{\\hfil$\\cdotp$\\hfil}}% \\copy0}} \\def\\|{\\ifmmode\\Vert\\else \\char`\\|\\fi} \\def\\sterling{{\\hbox{\\it\\char'44}}} \\def\\degrees{\\hbox{$^\\circ$}}% \\def\\degree{\\degrees}% \\def\\real{\\mathop{\\rm Re}\\nolimits}% \\def\\imag{\\mathop{\\rm Im}\\nolimits}% \\def\\tr{\\mathop{\\rm tr}\\nolimits}% \\def\\Tr{\\mathop{\\rm Tr}\\nolimits}% \\def\\Det{\\mathop{\\rm Det}\\nolimits}% \\def\\mod{\\mathop{\\rm mod}\\nolimits}% \\def\\wrt{\\mathop{\\rm wrt}\\nolimits}% \\def\\diag{\\mathop{\\rm diag}\\nolimits}% \\def\\TeV{{\\rm TeV}}% \\def\\GeV{{\\rm GeV}}% \\def\\MeV{{\\rm MeV}}% \\def\\keV{{\\rm keV}}% \\def\\eV{{\\rm eV}}% \\def\\Ry{{\\rm Ry}}% \\def\\mb{{\\rm mb}}% \\def\\mub{\\hbox{\\rm $\\mu$b}}% \\def\\nb{{\\rm nb}}% \\def\\pb{{\\rm pb}}% \\def\\fb{{\\rm fb}}% \\def\\cmsec{{\\rm cm^{-2}s^{-1}}}% \\def\\units#1{\\hbox{\\rm #1}} \\let\\unit=\\units \\def\\dimensions#1#2{\\hbox{$[\\hbox{\\rm #1}]^{#2}$}} \\def\\parenbar#1{{\\null\\! \\mathop{\\smash#1}\\limits ^{\\hbox{\\fiverm(--)}}% \\!\\null}}% \\def\\nunubar{\\parenbar{\\nu}} \\def\\ppbar{\\parenbar{p}} \\def\\buildchar#1#2#3{{\\null\\! \\mathop{\\vphantom{#1}\\smash#1}\\limits ^{#2}_{#3}% \\!\\null}}% \\def\\overcirc#1{\\buildchar{#1}{\\circ}{}} \\def\\sun{{\\hbox{$\\odot$}}}\\def\\earth{{\\hbox{$\\oplus$}}} \\def\\slashchar#1{\\setbox0=\\hbox{$#1$}% \\dimen0=\\wd0 \\setbox1=\\hbox{/} \\dimen1=\\wd1 \\ifdim\\dimen0>\\dimen1 \\rlap{\\hbox to \\dimen0{\\hfil/\\hfil}}% #1 \\else \\rlap{\\hbox to \\dimen1{\\hfil$#1$\\hfil}}% / \\fi}% \\def\\subrightarrow#1{% \\setbox0=\\hbox{% $\\displaystyle\\mathop{}% \\limits_{#1}$}% \\dimen0=\\wd0 \\advance \\dimen0 by .5em \\mathrel{% \\mathop{\\hbox to \\dimen0{\\rightarrowfill}}% \\limits_{#1}}}% \\newdimen\\vbigd@men \\def\\vbigl{\\mathopen\\vbig} \\def\\vbigm{\\mathrel\\vbig} \\def\\vbigr{\\mathclose\\vbig} \\def\\vbig#1#2{{\\vbigd@men=#2\\divide\\vbigd@men by 2 \\hbox{$\\left#1\\vbox to \\vbigd@men{}\\right.\\n@space$}}} \\def\\Leftcases#1{\\smash{\\vbigl\\{{#1}}} \\def\\Rightcases#1{\\smash{\\vbigr\\}{#1}}} \\def\\doublecolumns{\\relax}\\def\\enddoublecolumns{\\relax} \\def\\leftcolrule{\\relax}\\def\\rightcolrule{\\relax} \\def\\longequation{\\relax}\\def\\endlongequation{\\relax} \\def\\newcolumn{\\relax} \\def\\widetopinsert{\\topinsert}\\def\\widepageinsert{\\pageinsert} \\def\\forceleft{\\relax}\\def\\forceright{\\relax} \\def\\SetDoubleColumns#1{% \\imsg{The double column macros are not a part of mTeXsis.} \\imsg{If you want to use double column mode, get TXSdcol.tex} \\imsg{and add \\string\\input\\space TXSdcol.tex to your .tex file.} } \\def\\addTOC#1#2#3{\\relax}\\def\\Contents{\\relax} % \\newif\\ifContents % \\def\\ContentsSwitchtrue{\\Contentstrue}\\def\\ContentsSwitchfalse{\\Contentsfalse} \\def\\obsolete#1#2{\\let#1=#2\\relax #2} % \\let\\Input=\\input % \\newdimen\\colwidth \\colwidth=\\hsize % \\def\\ORGANIZATION{}% \\newhelp\\@utohelp{% loadstyle: The definition of the macro named above is actually contained^^J% in a style file, and so it cannot be used with mTeXsis. If you really^^J% need to load the definition from that file, you should do so explicitly^^J% at the begining of your manuscript file, with % '\\string\\input\\space stylefilename.txs'^^J} \\Ignore \\def\\loadstyle#1#2{% \\newlinechar=10 % \\errhelp=\\@utohelp % \\emsg{> Whoops! Trying to load \\string#1\\space from style file #2.}% \\errmessage{You cannot use macro definitions from style files in mTeXsis}} \\endIgnore \\hbadness=10000 % \\overfullrule=0pt % \\vbadness=10000 % \\ATunlock \\SetDate % \\ReadAUX % \\def\\fmtname{TeXsis}\\def\\fmtversion{2.17}% \\def\\revdate{1 January 1998}% \\def\\imsg#1{\\emsg{\\@comment #1}}% \\imsg{=========================================================== \\@comment} \\imsg{This is mTeXsis, the core macros from TeXsis.} \\imsg{You can get the complete TeXsis package (and avoid this annoying} \\imsg{advertisement) from ftp://lifshitz.ph.utexas.edu/texsis, } \\imsg{or from a CTAN server near you (in macros/texsis).} \\imsg{See the README and INSTALL files there for more information.} \\imsg{============================================================ \\@comment} \\emsg{m\\fmtname\\space version \\fmtversion\\space (\\revdate) loaded.}% \\ATlock % \\texsis % \\quoteoff \\global\\mathchardef\\DELTA=\"7001 \\global\\mathchardef\\LAMBDA=\"7003 \\global\\mathchardef\\XI=\"7004 \\global\\mathchardef\\SIGMA=\"7006 \\global\\mathchardef\\UPSILON=\"7007 \\global\\mathchardef\\OMEGA=\"700A \\global\\mathchardef\\Delta=\"7101 \\global\\mathchardef\\Lambda=\"7103 \\global\\mathchardef\\Xi=\"7104 \\global\\mathchardef\\Sigma=\"7106 \\global\\mathchardef\\Upsilon=\"7107 \\global\\mathchardef\\Omega=\"710A \\catcode`\\@=11 \\font\\tenmsa=msam10 \\font\\sevenmsa=msam7 \\font\\fivemsa=msam5 \\font\\tenmsb=msbm10 \\font\\sevenmsb=msbm7 \\font\\fivemsb=msbm5 \\newfam\\msafam \\newfam\\msbfam \\textfont\\msafam=\\tenmsa \\scriptfont\\msafam=\\sevenmsa \\scriptscriptfont\\msafam=\\fivemsa \\textfont\\msbfam=\\tenmsb \\scriptfont\\msbfam=\\sevenmsb \\scriptscriptfont\\msbfam=\\fivemsb \\def\\hexnumber@#1{\\ifnum#1<10 \\number#1\\else \\ifnum#1=10 A\\else\\ifnum#1=11 B\\else\\ifnum#1=12 C\\else \\ifnum#1=13 D\\else\\ifnum#1=14 E\\else\\ifnum#1=15 F\\fi\\fi\\fi\\fi\\fi\\fi\\fi} \\def\\msa@{\\hexnumber@\\msafam} \\def\\msb@{\\hexnumber@\\msbfam} \\global\\mathchardef\\boxdot=\"2\\msa@00 \\global\\mathchardef\\boxplus=\"2\\msa@01 \\global\\mathchardef\\boxtimes=\"2\\msa@02 \\global\\mathchardef\\square=\"0\\msa@03 \\global\\mathchardef\\blacksquare=\"0\\msa@04 \\global\\mathchardef\\centerdot=\"2\\msa@05 \\global\\mathchardef\\lozenge=\"0\\msa@06 \\global\\mathchardef\\blacklozenge=\"0\\msa@07 \\global\\mathchardef\\circlearrowright=\"3\\msa@08 \\global\\mathchardef\\circlearrowleft=\"3\\msa@09 \\global\\mathchardef\\rightleftharpoons=\"3\\msa@0A \\global\\mathchardef\\leftrightharpoons=\"3\\msa@0B \\global\\mathchardef\\boxminus=\"2\\msa@0C \\global\\mathchardef\\Vdash=\"3\\msa@0D \\global\\mathchardef\\Vvdash=\"3\\msa@0E \\global\\mathchardef\\vDash=\"3\\msa@0F \\global\\mathchardef\\twoheadrightarrow=\"3\\msa@10 \\global\\mathchardef\\twoheadleftarrow=\"3\\msa@11 \\global\\mathchardef\\leftleftarrows=\"3\\msa@12 \\global\\mathchardef\\rightrightarrows=\"3\\msa@13 \\global\\mathchardef\\upuparrows=\"3\\msa@14 \\global\\mathchardef\\downdownarrows=\"3\\msa@15 \\global\\mathchardef\\upharpoonright=\"3\\msa@16 \\let\\restriction=\\upharpoonright \\global\\mathchardef\\downharpoonright=\"3\\msa@17 \\global\\mathchardef\\upharpoonleft=\"3\\msa@18 \\global\\mathchardef\\downharpoonleft=\"3\\msa@19 \\global\\mathchardef\\rightarrowtail=\"3\\msa@1A \\global\\mathchardef\\leftarrowtail=\"3\\msa@1B \\global\\mathchardef\\leftrightarrows=\"3\\msa@1C \\global\\mathchardef\\rightleftarrows=\"3\\msa@1D \\global\\mathchardef\\Lsh=\"3\\msa@1E \\global\\mathchardef\\Rsh=\"3\\msa@1F \\global\\mathchardef\\rightsquigarrow=\"3\\msa@20 \\global\\mathchardef\\leftrightsquigarrow=\"3\\msa@21 \\global\\mathchardef\\looparrowleft=\"3\\msa@22 \\global\\mathchardef\\looparrowright=\"3\\msa@23 \\global\\mathchardef\\circeq=\"3\\msa@24 \\global\\mathchardef\\succsim=\"3\\msa@25 \\global\\mathchardef\\gtrsim=\"3\\msa@26 \\global\\mathchardef\\gtrapprox=\"3\\msa@27 \\global\\mathchardef\\multimap=\"3\\msa@28 \\global\\mathchardef\\therefore=\"3\\msa@29 \\global\\mathchardef\\because=\"3\\msa@2A \\global\\mathchardef\\doteqdot=\"3\\msa@2B \\let\\Doteq=\\doteqdot \\global\\mathchardef\\triangleq=\"3\\msa@2C \\global\\mathchardef\\precsim=\"3\\msa@2D \\global\\mathchardef\\lesssim=\"3\\msa@2E \\global\\mathchardef\\lessapprox=\"3\\msa@2F \\global\\mathchardef\\eqslantless=\"3\\msa@30 \\global\\mathchardef\\eqslantgtr=\"3\\msa@31 \\global\\mathchardef\\curlyeqprec=\"3\\msa@32 \\global\\mathchardef\\curlyeqsucc=\"3\\msa@33 \\global\\mathchardef\\preccurlyeq=\"3\\msa@34 \\global\\mathchardef\\leqq=\"3\\msa@35 \\global\\mathchardef\\leqslant=\"3\\msa@36 \\global\\mathchardef\\lessgtr=\"3\\msa@37 \\global\\mathchardef\\backprime=\"0\\msa@38 \\global\\mathchardef\\risingdotseq=\"3\\msa@3A \\global\\mathchardef\\fallingdotseq=\"3\\msa@3B \\global\\mathchardef\\succcurlyeq=\"3\\msa@3C 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\\global\\mathchardef\\shortmid=\"3\\msb@70 \\global\\mathchardef\\shortparallel=\"3\\msb@71 \\global\\mathchardef\\smallsetminus=\"2\\msb@72 \\global\\mathchardef\\thicksim=\"3\\msb@73 \\global\\mathchardef\\thickapprox=\"3\\msb@74 \\global\\mathchardef\\approxeq=\"3\\msb@75 \\global\\mathchardef\\succapprox=\"3\\msb@76 \\global\\mathchardef\\precapprox=\"3\\msb@77 \\global\\mathchardef\\curvearrowleft=\"3\\msb@78 \\global\\mathchardef\\curvearrowright=\"3\\msb@79 \\global\\mathchardef\\digamma=\"0\\msb@7A \\global\\mathchardef\\varkappa=\"0\\msb@7B \\global\\mathchardef\\hslash=\"0\\msb@7D \\global\\mathchardef\\hbar=\"0\\msb@7E \\global\\mathchardef\\backepsilon=\"3\\msb@7F \\def\\Bbb{\\ifmmode\\let\\next\\Bbb@\\else \\def\\next{\\errmessage{Use \\string\\Bbb\\space only in math mode}}\\fi\\next} \\def\\Bbb@#1{{\\Bbb@@{#1}}} \\def\\Bbb@@#1{\\fam\\msbfam#1} \\catcode`\\@=12 \\ATunlock % \\def\\refFormat{\\catcode`\\^^M=10} \\def\\Refs#1#2{Refs.~\\use{Ref.#1},\\use{Ref.#2}} \\def\\Refsand#1#2{Refs.~\\use{Ref.#1} and~\\use{Ref.#2}} \\def\\Refsrange#1#2{Refs.~\\use{Ref.#1}--\\use{Ref.#2}} \\superrefsfalse % \\def\\Chapter#1{Chapter~\\use{Chap.#1}} \\def ", "introduction": " ", "conclusions": "" }, "1002/1002.3394_arXiv.txt": { "abstract": "We discuss how the effective radius \\re\\ function (ERF) recently worked out by Bernardi et al. (2009) represents a new testbed to improve the current understanding of Semi-analytic Models of Galaxy formation. In particular, we here show that a detailed hierarchical model of structure formation can broadly reproduce the correct peak in the size distribution of local early-type galaxies, although it significantly overpredicts the number of very compact and very large galaxies. This in turn is reflected in the predicted size-mass relation, much flatter than the observed one, due to too large ($\\gtrsim 3$ kpc) low-mass galaxies ($<10^{11}$\\msun), and to a non-negligible fraction of compact ($\\lesssim 0.5-1$ kpc) and massive galaxies ($\\gtrsim 10^{11}$\\msun). We also find that the latter discrepancy is smaller than previously claimed, and limited to only ultracompact (\\re$\\lesssim 0.5$ kpc) galaxies when considering elliptical-dominated samples. We explore several causes behind these effects. We conclude that the former problem might be linked to the initial conditions, given that large and low-mass galaxies are present at all epochs in the model. The survival of compact and massive galaxies might instead be linked to their very old ages and peculiar merger histories. Overall, knowledge of the galactic stellar mass {\\em and} size distributions allows a better understanding of where and how to improve models. ", "introduction": "\\label{sec|intro} The formation and evolution of early-type galaxies, characterized by a dominant central stellar bulge component, is still a matter of debate. The seminal paper by \\citet[][]{Eggen62}, postulated that stars are formed in a single burst of star formation from gas falling towards the center, and the evolution is passive thereafter. Such a simple scenario might be difficult to reconcile with the standard cosmological paradigm of structure formation, in which dark matter halos grow hierarchically through merging. The most advanced and up-to-date Semi-analytical models (SAMs) of galaxy formation \\citep[e.g.,][]{Cole00,Benson03,Granato04,Granato06,Menci04,Ciras05,Khochfar05,Vittorini05,Bower06,Cattaneo06,Croton06,DeLucia06,Hopkins06,Lapi06,Shankar06,Monaco07,Somerville08SAM,Cook09,Fontanot09} still do not completely agree on the type of evolution undergone by massive galaxies (see also \\citealt{Dekel09}), on the fraction of stellar mass formed in the initial, gas-rich burst of star formation, and on the role played by the late evolution driven by major and minor mergers. Nevertheless, all models agree that galaxies must have been much more compact at the epoch of formation, owing to a denser universe, larger gas fractions in the progenitors, and more dissipation. The latter prediction has been confirmed by a number of deep observations of high redshift galaxies \\citep[e.g.,][]{Trujillo06,Vandokkum08,Cimatti08,Saracco08}, which have independently found early-type, high-redshift, massive galaxies to be a factor of a few more compact than local counterparts of the same stellar mass. Note, however, that several observational biases might limit the quality and reliability of some of these measurements \\citep[e.g.,][]{HopkinsRez,Mancini09,vandokkum09ALL}. It is still debated how these compact galaxies have evolved from high-redshifts increasing their sizes in a way to fall on the size-mass relation we observe today. As already extensively discussed by, e.g., \\citet{Shen03}, \\citet{ShankarBernardi}, \\PapI, the scatter around the local, median \\re-\\mstar\\ relation decreases with increasing mass and, at high stellar masses, is nearly independent of the age of the galaxies, a challenge for most galaxy evolution models. In particular, older galaxies are observed to have a steeper size-mass relation than younger systems. One possible model put forward to explain the strong size evolution of the red, massive high-$z$ galaxies is a sequence of minor, dry mergers. Such dynamical events can ``puff-up'' galaxies by adding mass in their outskirts, efficiently increasing their sizes, although limiting the growth of the stellar mass within a factor of $\\sim 2$ \\citep[e.g.,][]{KochfarSilk06origin,CiottiReview,Bernardi09,Bezanson09,Cimatti09,Hopkins09CompactGalx,Naab09,vanderwel09}. Recent numerical simulations have however shed doubts on the actual efficiency of mergers in significantly puffing-up compact galaxies, and coherently bringing them along the rather tight structural relations observed in the local Universe \\citep[][]{Nipoti09}. Moreover, hierarchical models suffer from the serious problem of failing in fully reproducing the local size-luminosity relation (see \\citealt{GonzalezSAM08}, \\PapInp, and references therein). Another model proposed in the recent Literature to increase the sizes of massive galaxies is by \\citet{Fan08}. They postulate that the evolution of compact galaxies undergoes a two-phase \\emph{expansion}: a first one caused by the sudden mass loss via quasar feedback, and a second one due to the slow mass loss via stellar winds during which the system slowly evolves towards a new equilibrium. Their model is broadly consistent with the data on the local size-mass relation. Given the large degree of freedom and significant uncertainties in current models of galaxy formation, it is necessary to look for other ways to test the validity of a given theory and/or discern ways to improve it. The aim of this paper is to provide such tests. We will show that the combined comparison with the size and mass distribution function of local galaxies can reveal interesting information on how to improve models of galaxy formation. In particular, in this work we will use a detailed hierarchical model of galaxy formation, show its failures and successes against available data, and discuss ways to improve it. We start in \\S~\\ref{sec|data} describing the data set we used. In \\S~\\ref{sec|models} we describe the hierarchical model adopted in this paper, and present its predicted size and mass distributions for spheroid-dominated galaxies. We will discuss in some detail the origin of the discrepancies between model predictions and the data, and possible ways to improve the model. We further discuss other aspects of the model in \\S~\\ref{sec|discu}, and conclude in \\S~\\ref{sec|conclu}. ", "conclusions": "\\label{sec|conclu} In this paper we make use of the data sets derived from SDSS DR6 by Bernardi et al. (2009) and \\citet{Hyde09a}, used to derive the size and stellar mass functions for a sample of early-type galaxies with concentration $C_r>2.86$, comprised of both ellipticals and S0 galaxies, and a sample dominated by ellipticals, respectively. We compare these statistical distributions with the hierarchical model by \\citet{Bower06}. The aim of this exercise is to show how the simultaneous comparison of the size and mass distributions can reveal interesting insights on how to improve the performance of theoretical models of galaxy evolution. We find, in agreement with previous studies, that this hierarchical model provides a poor match to the size-mass relation of local galaxies, irrespective of the exact sample we compare it with. In particular, the model tends to produce a much flatter relation than the one actually observed. This flattening is mainly produced by the combined effects of having, with respect to the local data, too large ($\\sim 3$ kpc) low-mass galaxies ($<10^{11}$\\msun), and of having a non-negligible fraction of compact galaxies ($\\lesssim 0.5-1$ kpc) at high masses ($\\gtrsim 10^{11}$\\msun). Such discrepancies are reflected in the predicted size distribution. Although the model produces a size distribution in broad agreement with the data, it tends to overproduce the number of large galaxies beyond the peak ($\\gtrsim 3$ kpc), and the number of very compact galaxies ($\\lesssim 1$ kpc). We discussed that the former issue is present at all epochs, and it might therefore be linked to how spheroids are formed in the first place, either from not properly treating initial disk instabilities and/or computing the sizes of remnants in gas-rich mergers. Regarding the overproduction of compact and massive (\\mstar$\\sim (0.5-1)\\times 10^{11}$\\msun) galaxies with respect to the data, already pointed out in the recent Literature, we find it to be less prominent than previously claimed, and confined to only ultracompact galaxies (\\re$\\lesssim 0.5$ kpc) when considering only ellipticals. We discuss two possible reasons behind the survival of such compact galaxies until the present epoch. First, we find that model early-type galaxies tend to be significantly older than those in SDSS. This in turn might induce more compact galaxies at fixed stellar mass, given that galaxies formed at higher redshifts are more compact (see \\PapInp). We also find that model early-type compact galaxies underwent peculiar merging histories characterized by extremely compact progenitors, that could prevent them to efficiently grow their sizes." }, "1002/1002.4782_arXiv.txt": { "abstract": "The prospects for future blazar surveys by next-generation very high energy (VHE) gamma-ray telescopes, such as Advanced Gamma-ray Imaging System (AGIS) and Cherenkov Telescope Array (CTA), are investigated using the latest model of blazar luminosity function and its evolution which is in good agreement with the flux and redshift distribution of observed blazars as well as the extragalactic gamma-ray background. We extend and improve the template of spectral energy distributions (SEDs) based on the blazar SED sequence paradigm, to make it reliable also in the VHE bands (above 100 GeV) by comparing with the existing VHE blazar data. Assuming the planned CTA sensitivities, a blind survey using a total survey time of $\\sim 100$ hrs could detect $\\sim 3$ VHE blazars, with larger expected numbers for wider/shallower surveys. We also discuss following-up of {\\it Fermi} blazars. Detectability of VHE blazars in the plane of {\\it Fermi} flux and redshift is presented, which would be useful for future survey planning. Prospects and strategies are discussed to constrain the extragalactic background light (EBL) by using the absorption feature of brightest blazar spectra, as well as cut-offs in the redshift distribution. We will be able to get useful constraints on EBL by VHE blazars at different redshifts ranging 0.3--1 TeV corresponding to $z=$0.10--0.36. ", "introduction": "\\label{sec:intro} Very high energy (VHE; above 100 GeV) gamma-ray astronomy has now firmly been established by the observations of the state-of-the-art imaging atmospheric Cherenkov Telescopes (IACTs) such as H.E.S.S., MAGIC, and VERITAS (see {de Angelis}, {Mansutti}, \\& {Persic} 2008; Mori 2009, for reviews). Further progress is anticipated in the near future by the planned next-generation IACTs such as Cherenkov Telescope Array (CTA) and Advanced Gamma-ray Imaging System (AGIS). The sensitivities of all-sky monitoring VHE gamma-ray experiments are also expected to improve by the future projects such as the High Altitude Water Cherenkov Experiment (HAWC) and the Tibet-III/MD experiment. Current IACTs have already found $\\sim$ 100 VHE sources including $\\sim$ 25 blazars. Blazars, a class of active galactic nuclei (AGNs), are the dominant population in the extragalactic gamma-ray sky. Almost all of the extragalactic sources detected by EGRET (Energetic Gamma-Ray Experiment Telescope) on board the Compton Gamma Ray Observatory are blazars ({Hartman} {et~al.} 1999). Moreover, 3 months bright source and 11 months catalog by the Fermi gamma-ray space telescope ({\\it Fermi}) have recently also showed that most of the extragalactic sources are blazars ({Abdo} {et~al.} 2009a, 2009c, 2010), and we expect that more than 1000 blazars will be detected by {\\it Fermi} in the near future (e.g. {Narumoto} \\& {Totani} 2006; {Dermer} 2007; {Inoue} \\& {Totani} 2009). The number of VHE blazars are expected to dramatically increase with the improved next-generation IACT sensitivity. Therefore, it would be possible to do a statistical study of VHE blazars in the CTA/AGIS era, which would provide a crucial key to understand AGN populations and high energy phenomena around super massive black holes in AGNs and jets. The purpose of this paper is to study the prospect of future blazar surveys by IACTs, especially for the statistical power of future VHE blazar sample that can be obtained by realistic observing time of next-generation IACTs. For this purpose, blazar gamma-ray luminosity function (GLF) and spectral energy distribution (SED) are needed. The blazar GLF has been studied in detail by many papers ({Padovani} {et~al.} 1993; {Stecker}, {Salamon}, \\& {Malkan} 1993; {Salamon} \\& {Stecker} 1994; {Chiang} {et~al.} 1995; {Stecker} \\& {Salamon} 1996; {Chiang} \\& {Mukherjee} 1998; {M{\\\"u}cke} \\& {Pohl} 2000; {Narumoto} \\& {Totani} 2006; {Dermer} 2007; {Inoue} \\& {Totani} 2009). {Inoue} \\& {Totani} (2009) (hereafter IT09) has recently presented a new blazar GLF taking into account the blazar SED sequence (see \\S \\ref{subsec:glf}), which is in nice agreement with the CGRO/EGRET and Fermi/LAT data. We utilize this IT09 model to predict the expected number and distributions of physical quantities of VHE blazars in future IACT surveys. Since the SED model of IT09 was constrained only at the photon energies under GeV, we construct a new blazar SED template by modifying that used in IT09 in accordance with the available VHE blazar data. By using our updated blazar sequence and GLF model, it is possible for us to make predictions for future VHE gamma-ray observations, which is the most reliable based on available observed data. Extragalactic background light (EBL) in the optical and infrared bands contains the information about the history of star formation activity in the universe, and knowing EBL quantitatively is an important step to understand galaxy formation in the cosmological context. However, it is hard to measure EBL spectrum directly, mainly because of the difficulty in subtracting foreground emission (see {Hauser} \\& {Dwek} 2001, for reviews). VHE observations provide a completely independent constraint on EBL, since VHE gamma-ray photons propagating the universe are absorbed via electron-positron pair creation with the EBL photons ({Gould} \\& {Schr{\\'e}der} 1966; {Jelley} 1966). Some useful limits have already been obtained by VHE blazar observations ({Aharonian} {et~al.} 2006a; {MAGIC Collaboration}, {Albert}, {et~al.} 2008), up to the redshift of $z=0.536$ by using 3C279 data. The next generation IACTs will shed further light on this issue, and we discuss the prospect about this as a particular application of our study. This paper is organized as follows. We introduce our updated blazar SED template and GLF model, as well as the model of VHE gamma-ray absorptions by EBL in \\S \\ref{sec:model}. In \\S \\ref{sec:cta}, we make predictions for the expected number and statistics of future VHE blazar surveys assuming some observing modes. We discuss the prospect for the determination of EBL by VHE blazars in \\S \\ref{sec:dis}. Summary is given in \\S \\ref{sec:sum}. Throughout this paper, we adopt the standard cosmological parameters of ($h,\\Omega_M,\\Omega_\\Lambda$)=(0.7,0.3,0.7). ", "conclusions": "\\label{sec:sum} In this paper, we estimated the expected source counts and redshift distribution of VHE blazars for the next generation IACTs such as CTA and AGIS missions based on the latest blazar GLF model of {Inoue} \\& {Totani} (2009). For this purpose, we developed a new SED sequence formula of blazars taking into account the latest VHE data, while the previous sequence formulae were constructed using only data at photon energies below the EGRET/Fermi energy band. The parameters of our blazar GLF are also refined with this new SED formula, by fitting to the GeV blazar data. Our modeling does not include time variabilities of VHE blazars, and blazars at flaring states would be more easily detected than estimated here. We made predictions for future VHE blazar survey in two observing modes: one is a blind survey in a blank field, and another is a following up survey of {\\it Fermi} blazars. We found that CTA will detect a few VHE blazars by a blind survey using a total survey time of 100 hours. Therefore a large amount of observing time ($\\gtrsim 1000$ hrs) is required to construct a statistically large sample of blazars selected only by VHE bands. However, this suggests that blazar contamination in the Galactic plane survey should not be significant even in the era of the next generation IACTs. We also found that future all sky gamma-ray detectors such as HAWC and Tibet-III/MD, will detect only a few VHE blazars in one year survey in the entire sky. The survey design for a follow-up survey of Fermi blazars should be dependent on the scientific purposes. Here we presented a plot for regions in the Fermi flux versus redshift plane where the Fermi blazars can be detected by VHE observations for several different sensitivities. As a particular example of Fermi blazar follow-up surveys, we considered a survey for the purpose of determination of EBL by VHE observation. CTA can observe VHE blazars that are sufficiently bright to get detailed spectra with high S/N in the redshift range of $z \\sim$ 0.10--0.36, corresponding to the absorption cut-off energy 1--0.3 TeV, and hence we can constrain not only EBL flux but also its spectra and/or redshift evolution. It will also be possible to construct a statistically large sample ($\\gtrsim 30$) of blazars at $z \\sim $ 0.10--0.36 to constrain EBL by the sharp break in the redshift distribution. This approach could avoid or minimize the uncertainty about intrinsic blazar spectra, and hence could be complementary to using a few spectra of brightest blazars. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. This work was supported by the Grant-in-Aid for the Global COE Program \"The Next Generation of Physics, Spun from Universality and Emergence\" from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan. YI acknowledges support by the Research Fellowship of the Japan Society for the Promotion of Science (JSPS). \\appendix" }, "1002/1002.1028_arXiv.txt": { "abstract": "DAMA/LIBRA is running at the Gran Sasso National Laboratory of the I.N.F.N.. Here the results obtained with a further exposure of 0.34 ton $\\times$ yr are presented. They refer to two further annual cycles collected one before and one after the first DAMA/LIBRA upgrade occurred on September/October 2008. The cumulative exposure with those previously released by the former DAMA/NaI and by DAMA/LIBRA is now 1.17 ton $\\times$ yr, corresponding to 13 annual cycles. The data further confirm the model independent evidence of the presence of Dark Matter (DM) particles in the galactic halo on the basis of the DM annual modulation signature (8.9 $\\sigma$ C.L. for the cumulative exposure). In particular, with the cumulative exposure the modulation amplitude of the {\\it single-hit} events in the (2 -- 6) keV energy interval measured in NaI(Tl) target is ($0.0116 \\pm 0.0013$) cpd/kg/keV; the measured phase is ($146 \\pm 7$) days and the measured period is ($0.999 \\pm 0.002$) yr, values well in agreement with those expected for the DM particles. ", "introduction": "The former DAMA/NaI \\cite{prop,allDM,Nim98,Sist,RNC,ijmd,ijma,epj06,ijma07,chan,wimpele,ldm,allRare,IDM96} and the present DAMA/LIBRA \\cite{modlibra,perflibra,papep} experiments at the Gran Sasso National Laboratory have the main aim to investigate the presence of Dark Matter particles in the galactic halo by exploiting the model independent Dark Matter annual modulation signature originally suggested in the mid 80's in ref. \\cite{Freese}. In fact, as a consequence of its annual revolution around the Sun, which is moving in the Galaxy travelling with respect to the Local Standard of Rest towards the star Vega near the constellation of Hercules, the Earth should be crossed by a larger flux of Dark Matter particles around $\\sim$ 2 June (when the Earth orbital velocity is summed to the one of the solar system with respect to the Galaxy) and by a smaller one around $\\sim$ 2 December (when the two velocities are subtracted). Thus, this signature has a different origin and peculiarities than the seasons on the Earth and than effects correlated with seasons (consider the expected value of the phase as well as the other requirements listed below). This annual modulation signature is very distinctive since the effect induced by DM particles must simultaneously satisfy all the following requirements: the rate must contain a component modulated according to a cosine function (1) with one year period (2) and a phase that peaks roughly around $\\simeq$ 2$^{nd}$ June (3); this modulation must only be found in a well-defined low energy range, where DM particle induced events can be present (4); it must apply only to those events in which just one detector of many actually ``fires'' ({\\it single-hit} events), since the DM particle multi-interaction probability is negligible (5); the modulation amplitude in the region of maximal sensitivity must be $\\lsim$7$\\%$ for usually adopted halo distributions (6), but it can be larger in case of some possible scenarios such as e.g. those in refs. \\cite{Wei01,Fre04}. This offers an efficient DM model independent signature, able to test a large interval of cross sections and of halo densities; moreover, the use of highly radiopure NaI(Tl) scintillators as target-detectors assures sensitivity to wide ranges of DM candidates, of interaction types and of astrophysical scenarios. It is worth noting that only systematic effects or side reactions able to simultaneously fulfil all the 6 requirements given above (and no one has ever been suggested) and to account for the whole observed modulation amplitude might mimic this DM signature. \\vspace{0.5cm} The DAMA/LIBRA set-up, whose description, radiopurity and main features are discussed in details in ref. \\cite{perflibra} has firstly been upgraded in September/October 2008: i) one detector has been recovered by replacing a broken PMT (see ref. \\cite{modlibra}); ii) a new optimization of some PMTs and HVs has been performed; iii) all the transient digitizers recording the shape of the pulse have been replaced with new ones, the U1063A Acqiris 8-bit 1GS/s DC270 High-Speed cPCI Digitizers; iv) a new DAQ with optical read-out has been installed. Also during this upgrade the operations involving the handling of the sensitive part of the setup and the shield have been performed in HP Nitrogen atmosphere. The upgrade has allowed to enlarge the sensitive mass and to improve general features. Here we just remind that the sensitive part of this set-up is made of 25 highly radiopure NaI(Tl) crystal scintillators (5-rows by 5-columns matrix) having 9.70 kg mass each one. In each detector two 10 cm long special quartz light guides act also as optical windows on the two end faces of the crystal and are coupled to two low background photomultipliers working in coincidence at single photoelectron level. The detectors are housed in a sealed low-radioactive copper box installed in the center of a low-radioactive Cu/Pb/Cd-foils/polyethylene/paraffin shield; moreover, about 1 m concrete (made from the Gran Sasso rock material) almost fully surrounds (mostly outside the barrack) this passive shield, acting as a further neutron moderator. A threefold-levels sealing system excludes the detectors from the environmental air of the underground laboratory \\cite{perflibra}. A hardware/software system to monitor the running conditions is operative and self-controlled computer processes automatically control several parameters and manage alarms. Moreover: i) the light response ranges from 5.5 to 7.5 photoelectrons/keV, depending on the detector; ii) the hardware threshold of each PMT is at single photoelectron (each detector is equipped with two low background photomultipliers working in coincidence); iii) energy calibration with X-rays/$\\gamma$ sources are regularly carried out down to few keV; iv) the software energy threshold of the experiment is 2 keV; v) both {\\it single-hit} events (where just one of the detectors fires) and {\\it multiple-hit} events (where more than one detector fires) are acquired; v) the data are collected up to the MeV region despite the optimization is performed for the lower one. For the radiopurity, the procedures and further details see ref. \\cite{modlibra,perflibra}. The data of the former DAMA/NaI (0.29 ton $\\times$ yr) and those of the first 4 annual cycles of DAMA/LIBRA (total exposure 0.53 ton$\\times$yr) have already given positive model independent evidence for the presence of DM particles in the galactic halo with high confidence level on the basis of the DM annual modulation signature \\cite{modlibra}. In this paper the model independent results with other two annual cycles DAMA/LIBRA-5,6 are presented. As mentioned, the data of the first cycle have been collected in the same conditions as DAMA/LIBRA-1,2,3,4 \\cite{modlibra,perflibra}, while the data of DAMA/LIBRA-6 have been taken after the above mentioned 2008 upgrade. ", "conclusions": "The new annual cycles DAMA/LIBRA-5,6 have further confirmed a peculiar annual modulation of the {\\it single-hit} events in the (2--6) keV energy region satisfying the many requests of the DM annual modulation signature; the total exposure by former DAMA/NaI and present DAMA/LIBRA is 1.17 ton $\\times$ yr. In fact, as required by the DM annual modulation signature: 1) the {\\it single-hit} events show a clear cosine-like modulation as expected for the DM signal; 2) the measured period is equal to $(0.999\\pm 0.002)$ yr well compatible with the 1 yr period as expected for the DM signal; 3) the measured phase $(146\\pm7)$ days is well compatible with the roughly $\\simeq$ 152.5 days expected for the DM signal; 4) the modulation is present only in the low energy (2--6) keV energy interval and not in other higher energy regions, consistently with expectation for the DM signal; 5) the modulation is present only in the {\\it single-hit} events, while it is absent in the {\\it multiple-hit} ones as expected for the DM signal; 6) the measured modulation amplitude in NaI(Tl) of the {\\it single-hit} events in the (2--6) keV energy interval is: $(0.0116 \\pm 0.0013)$ cpd/kg/keV (8.9 $\\sigma$ C.L.). No systematic or side processes able to simultaneously satisfy all the many peculiarities of the signature and to account for the whole measured modulation amplitude is available. Further work is in progress.\\\\" }, "1002/1002.0504_arXiv.txt": { "abstract": "We report the discovery of two new Milky Way satellites in the neighboring constellations of Pisces and Pegasus identified in data from the Sloan Digital Sky Survey. Pisces II, an ultra-faint dwarf galaxy lies at the distance of $\\sim 180$ kpc, some 15$^\\circ$ away from the recently detected Pisces I. Segue 3, an ultra-faint star cluster lies at the distance of 16 kpc. We use deep follow-up imaging obtained with the 4-m Mayall telescope at Kitt Peak National Observatory to derive their structural parameters. Pisces II has a half-light radius of $\\sim$60 pc, while Segue 3 is twenty times smaller at only 3pc. ", "introduction": "Ultra-faint satellites of the Milky Way include dwarf galaxies ~\\citep{Wi05,Zu06a,Zu06b,Be06a,Be07,Ir07,Be08} and star clusters ~\\citep{Ko07}, as well as objects with intermediate properties ~\\citep{Wal07,Be09, Ni09}. Defined by their extremely low surface brightness, these systems could only be detected with a massive multi-band imaging campaign like the Sloan Digital Sky Survey (SDSS). In this {\\it Letter}, we announce the discovery of two further Milky Way satellites in the Southern Galactic portion of the SDSS SEGUE survey. They lie in adjacent constellations and each extend only a couple of arc-minutes on the sky. However, their heliocentric distances differ by an order of magnitude, and, hence, so do their physical sizes. We name the dwarf galaxy in the constellation of Pisces, lying at the heliocentric distance of $\\sim 180$ kpc and measuring $\\sim 120$ pc across, Pisces II. This is the second Galactic stellar halo sub-structure in Pisces - the first, Pisces I was announced earlier this year by ~\\citet{Wat09}. Pisces I is much closer and more dispersed on the sky: it is at least several degrees across and lies at $\\sim 80$ kpc. Our second discovery is a feeble cluster of stars in the constellation of Pegasus. It has a half-light radius of 3 pc and lies at a heliocentric distance of 16 kpc. We name it Segue 3, after SEGUE, the imaging survey in the data of which it was found. In the analysis presented in this Letter we have extinction-corrected all magnitudes using the maps of \\citet{Sc98}. \\begin{figure*}[t] \\begin{center} \\includegraphics[width=0.9\\textwidth]{segue3_pisces2_figure1_sdss.ps} \\caption{The SDSS view of Pisces II (upper row) and Segue 3 (lower row): {\\it Left:} Density of stars in the SDSS catalogue centered on the object. The stars in the $10^\\prime \\times 10^\\prime$ area are binned into $15\\times15$ bins and smoothed with a Gaussian with FWHM of 1.5 pixel. {\\it Middle Left:} CMD of all stars in a circle of radius $1\\farcm2$ (marked on the left panel), dominated by the satellite's members. {\\it Middle Right:} Comparison CMD of stars within the annulus $7^\\prime$ to $7\\farcm1$ showing the foreground. {\\it Right:} Difference in Hess Diagrams. Pisces II populations (top right) can be hinted at by the red giant branch and blue horizontal branch. Segue 3 (bottom right) shows an obvious main sequence. Ridge-lines of M92 (red, [Fe/H]$\\sim$-2.3) and M13 (blue, [Fe/H]$\\sim$-1.55) are over-plotted. For Segue 3, we also show a mask built using M92 ridge-line to select possible red giant stars for luminosity calculation (lower middle left panel).} \\label{fig:fig1_sdss} \\end{center} \\end{figure*} \\begin{deluxetable}{lcc} \\tablecaption{Properties of Pisces II and Segue 3 \\label{tbl:pars}} \\tablewidth{0pt} \\tablehead{ \\colhead{Parameter} & {Pisces II} & {Segue 3}} \\startdata RA (J2000) & $22:58:31 \\pm 6$ & $21:21:31 \\pm 4$ \\\\ Dec (J2000) & $+05:57:09 \\pm 4$ & $+19:07:02 \\pm 4$ \\\\ Galactic $\\ell$ & $79.21^\\circ$ & $69.4^\\circ$ \\\\ Galactic $b$ & $-47.11^\\circ$ & $-21.27^\\circ$ \\\\ $r_h$ (Plummer) & $1\\farcm1 \\pm 0\\farcm1$ & $0\\farcm65 \\pm 0\\farcm1$ \\\\ $\\theta$ & $77^\\circ \\pm 12^\\circ$ & $215^\\circ \\pm 20^\\circ$\\\\ $e$ & $0.4 \\pm 0.1$ & $0.3 \\pm 0.2$\\\\ (m$-$M)$_0$ & $21\\fm3$ & $16\\fm1$ \\\\ M$_{\\rm tot,V}$ & $-5\\fm0$ & $-1\\fm2$ \\enddata \\tablenotetext{*}{Magnitudes are accurate to $\\sim \\pm 0\\fm5$ and are corrected for the Galactic foreground reddening.} \\label{tab:struct} \\end{deluxetable} ", "conclusions": "Pisces~II is close on the sky to Pisces~I, discovered as an overdensity of RR Lyraes by \\citet{Wat09} and confirmed spectroscopically by \\citet{Ko09}. But, at a heliocentric distance of $\\sim 180$ kpc, Pisces~II is almost twice as far as away. The extent of Pisces~I is probably considerable, at least as judged from the RR Lyrae populations (see Figure 12 of Watkins et al. 2009). This might imply the break-up of a substantial satellite galaxy moving on a radial orbit, in which case it is natural to interpret Pisces~II as a further fragment or companion. Pisces~II is very similar in morphology, size and luminosity to a number of recent discoveries such as Hercules, Leo~IV and Leo~V~\\citep{Be07, Be08}. They also all lie at similar heliocentric distances of $\\sim 150$ kpc. All four have extended BHB populations enshrouding them. Segue~3 is a very close relative of Koposov~1 and 2, the ultrafaint star clusters at distances of $\\sim 50$ kpc~\\citep{Ko07}. All three objects have a similar size ($\\sim 3$ pc), luminosity $M_{\\rm V}\\sim -1$ and contain only a few tens of stars. Unlike Koposov~1 and 2, Segue~3 is much closer, and might even be a part of the Hercules-Aquila Cloud~\\citep{Be07b,Wat09}. The evolution of such objects is known to proceed with pronounced mass segregation. The very few heavy stars sink to the centre, and the lighter stars are ejected to form a diffuse corona. It would be interesting to verify this prediction with spectrocopic surveys of the objects." }, "1002/1002.0397_arXiv.txt": { "abstract": "% The high-resolution pictures of the solar photosphere from space based 50 cm Solar Optical Telescope (SOT) onboard {\\it Hinode} spacecraft, are now routinely observed. Such images of a $\\delta$-sunspot in NOAA 10930 were obtained by {\\it Hinode} during 13 December 2006 while a X-class flare occurred in this active region. Two bright ribbons were visible even in white light and G-band images apart from chromospheric Ca II H images. We register the sunspot globally using cross-correlation technique and analyse local effects during flare interval. We find that during flare the penumbral filaments show lateral motion. Also, we locate two patches, one in either polarity, which show converging motion towards the polarity inversion line (PIL). In Ca II H images we find kernel with pre-flare brightening which lie along the PIL. ", "introduction": "% The flares in the solar corona are believed to be due to sudden restructuring of the stressed magnetic field. The energy is released in the form of thermal as well as non-thermal radiation and energetic charged particles. In powerful flares the energetic particles can penetrate the dense chromosphere to reach down to the photosphere where they heat-up the photosphere leading to white-light flares. It has been observed that photospheric changes are accompanied during these highly energetic events in the form of: (i) change in morphology, (ii) change in magnetic flux, (iii) change in magnetic shear angle (the angle between observed field azimuth and potential field azimuth) and (iv) proper motion. Here we focus on the local changes, i.e., changes seen in small-scale features like penumbral filaments during flares. Such studies require seeing-free high-resolution observations at a high-cadence. This is possible with the 50 cm Solar Optical Telescope (SOT) onboard {\\it Hinode} spacecraft \\citep{Kosugi2007,Tsuneta2008,Ichimoto2008,Suematsu2008}. Here, we present the observations of a $\\delta$-sunspot in active region NOAA 10930 during a X-class flare on 13 December 2006 at 02:20 UT by {\\it Hinode}. The two ribbons could be seen in G-band and Fe I 630.2 nm Stokes-I and V images \\citep{Isobe2007}. Earlier, we had reported the lateral motion of penumbral filaments during the flare interval \\citep{Gosain2009}. Here, we present the converging motion of the two patches, one in either polarity, located on either side of the PIL. Also, the pre-flare brightening in kernels located along the PIL as seen in Ca II H line, is discussed. ", "conclusions": "We have studied the evolution of small scale features in a flaring sunspot using high resolution space based observations. The fine structure of the sunspot penumbra as seen in high resolution G-band images is believed to outline the magnetic field lines. Using the penumbral filaments as a proxy for magnetic field structure of the sunspot we study the changes in the structure during X-class flare of 13 December 2006. A two phased motion is seen in patches of either poalrity, i.e., boxes `1' and `2'. First phase of motion lasts for about four minutes, directed away from the neutral line, and another phase of motion lasts for more than 40 minutes, directed towards neutral line, i.e., converging. Further, we notice that in Ca II H images the pre-flare brightening is clearly visible in elongated kernels, located along the PIL. This brightening is seen as early as 16 minutes prior to the flare onset. To understand the changes in magnetic field configuration during flares one requires high-cadence vector magnetograms obtained with high-resolution which are expected from upcoming space missions like Helioseismic and Magnetic Imager (HMI) and Solar Orbiter." }, "1002/1002.3408.txt": { "abstract": "Using numerical ray tracing, the paper studies how the average distance modulus in an inhomogeneous universe differs from its homogeneous counterpart. The averaging is over all directions from a fixed observer not over all possible observers (cosmic), thus it is more directly applicable to our observations. Unlike previous studies, the averaging is exact, non-perturbative, and includes all possible non-linear effects. The inhomogeneous universes are represented by Sweese-cheese models containing random and simple cubic lattices of mass-compensated voids. The Earth observer is in the homogeneous cheese which has an Einstein - de Sitter metric. For the first time, the averaging is widened to include the supernovas inside the voids by assuming the probability for supernova emission from any comoving volume is proportional to the rest mass in it. For voids aligned in a certain direction, there is a cumulative gravitational lensing correction to the distance modulus that increases with redshift. That correction is present even for small voids and depends on the density contrast of the voids, not on their radius. Averaging over all directions destroys the cumulative correction even in a non-randomized simple cubic lattice of voids. Despite the well known argument for photon flux conservation, the average distance modulus correction at low redshifts is not zero due to the peculiar velocities. A formula for the maximum possible average correction as a function of redshift is derived and shown to be in excellent agreement with the numerical results. The formula applies to voids of any size that: (1) have approximately constant densities in their interior and walls, (2) are not in a deep nonlinear regime. The actual average correction calculated in random and simple cubic void lattices is severely damped below the predicted maximum. That is traced to cancelations between the corrections coming from the fronts and backs of different voids at the same redshift from the observer. The calculated correction at low redshifts allows one to readily predict the redshift at which the averaged fluctuation in the Hubble diagram is below a required precision and suggests a method to extract the background Hubble constant from low redshift data without the need to correct for peculiar velocities. ", "introduction": "Every reasonable model of the universe is matter dominated in the past while the structure is being formed. Adding a cosmological constant has little impact \\cite{bolejko} on the development up until recent times when it starts dominating the matter. That motivates choosing the spatially-flat matter-only Einstein-de Sitter (EdS) model for the homogeneous regions of the Swiss-cheese. The matter includes both visible and dark varieties. This is a convenient playground to explore the question whether voids in the matter distribution can mimic the effect of dark energy. The Hubble constant of the model is $H_0 = 70 \\, $ km/s/Mpc. The big bang time is conventionally chosen as $t_{bb}=0$, the current age of the model is $t_0=2/(3H_0)=2857$ Mpc (9312 Myr) and the scale factor as a function of cosmic time is $a(t)=(t/t_0)^{2/3}$. The matter density today is $\\bar{\\rho}(t_0)=3H_0^2/(8\\pi G)$ and as a function of time is: \\begin{equation}\\label{rhot} \\bar{\\rho}(t)=\\frac{\\bar{\\rho}(t_0)}{a(t)^3} = \\frac{1}{6\\pi G \\, t^2}\\,. \\end{equation} ", "conclusions": "" }, "1002/1002.2392_arXiv.txt": { "abstract": "The origins of irregular satellites of the giant planets are an important piece of the giant ``puzzle'' that is the theory of Solar System formation. \\corr[1]{It is well established that they are not \\textit{in situ} formation objects, around the planet, as are believed to be the regular ones. Then, the most plausible hypothesis to explain their origins is that they formed elsewhere and were captured by the planet. However, captures under restricted three-body problem dynamics have temporary feature, which makes necessary the action of an auxiliary capture mechanism. Nevertheless, there not exist one well established capture mechanism}.% \\corr[NEW]{In this work, we tried to understand which aspects of a binary-asteroid capture mechanism could favour the permanent capture of one member of a binary asteroid.}% We performed more than eight thousand numerical simulations of capture trajectories considering the four-body dynamical system Sun, Jupiter, Binary-asteroid. We restricted the problem to the circular planar prograde case, and time of integration to $10^4$ years. % With respect to the binary features, we noted that 1) tighter binaries are much more susceptible to produce permanent captures than the large separation-ones. We also found that 2) the permanent capture probability of the minor member of the binary is much more expressive than the major body permanent capture probability. % On the other hand, among the aspects of capture-disruption process, 4) a pseudo eastern-quadrature was noted to be a very likely capture angular configuration at the instant of binary disruptions. In addition, we also found that the 5) capture probability is higher for binary asteroids which disrupt in an inferior-conjunction with Jupiter. % These results show that the Sun plays a very important role on the capture dynamic of binary asteroids. ", "introduction": "\\gram{The existence of more than 350 natural satellites is known}, from which approximately 50\\,\\% are planetary ones. \\corr[34]{An interesting point about this number, is that before 1997 just a tenth of such objects was known, i.e., the new ``CCD observational era'' allowed this number to increase by an order of magnitude within just a half decade \\citep{gladmanetal98, gladmanetal00, gladmanetal01, sheppardjewitt03, holmanetal04, kavelaarsetal04, sheppardetal05, sheppardetal06}}. The planetary satellites can be distinguished into two characteristic groups: regulars and irregulars \\citep{kuiper56, peale99}. The first group, is characterized by small values of semi-major axis, eccentricities and inclinations. These characteristics are a strong signature of \\textit{in situ}-formation through matter accretion from the circumplanetary disc \\citep{ luninestevenson82, vieiranetowinter01, canupward02, canupward06, mosqueiraestrada03, sheppardjewitt03}. In \\gram{contrast}, the irregular satellites have large values of semi-major axis \\citep{burns86}, \\gram{often} high eccentricities and inclinations. A large part of irregular satellites have \\gram{retrograde} orbital inclinations higher than 90 degrees \\citep{jewitthaghighipour07}. Another important characteristic of the irregular ones are the family groups, i.e., satellite groups characterized by similar orbital elements \\corr[5]{\\citep{gladmanetal01, kavelaarsetal04}}. % These peculiar characteristics are incompatible with the \\textit{in situ}-formation model through matter accretion \\citep{kuiper56}, and since they are the majority group of planetary satellites in the solar system, there exists a large scientific interest about their origin. Then, the most plausible hypothesis to explain their origins is that they formed elsewhere and were captured by the planet \\corr[6]{\\citep{kuiper56, heppenheimerporco77, pollacketal79, colombofranklin71}}. However, \\gram{many studies} have shown that gravitational captures under three-body-dynamics are temporary \\citep{everhart73, heppenheimerporco77, carusivalsecchi79, bennermckinnon95, vieiranetowinter01, wintervieiraneto01}. This fact has induced researchers to propose some auxiliary capture mechanism. Among others, we point out four mostly well known: \\begin{enumerate} \\item Gas drag capture \\citep{pollacketal79, cukburns04}: A \\gram{temporarily} captured asteroid becomes permanently captured through kinetic energy decrease due to gas drag inside the circumplanetary disk of gas and dust; \\item Pull-Down capture \\citep{heppenheimerporco77, vieiranetoetal04, oliveiraetal07}: A temporary captured asteroid becomes permanently captured due to an increase of the Hill's radius of the planet. This increase in Hill's radius occurs due to either planet's mass growth or planet's migration away from the Sun; \\item \\gram{Close-approach} interaction captures \\citep{colombofranklin71, tsui99, tsui00, astakhovetal03, nesvornyetal03, funatoetal04}: A \\gram{temporarily} captured asteroid becomes permanently captured through energy and angular momentum exchanges with an existing satellite; \\item Capture of binary-asteroids \\citep{agnorhamilton06, vokrouhlickyetal08}: One member of a binary-asteroid becomes permanently captured when the binary approaches the planet and disrupts. \\end{enumerate} The capture mechanism of binary-asteroids is very interesting since the present observations have shown an increasing number of such systems in the main populations of such objects as the Kuiper Belt, Main Belt and Near Earth Asteroids \\citep{noll06}. \\citet{agnorhamilton06} presented numerical simulations of close encounters between Neptune and a binary-asteroid, where they considered \\gram{one asteroid comparable} to Triton and \\corr{a secondary}, with equal mass or one order of magnitude \\corr{lower}. Their results show that is possible to \\gram{disrupt} the binary when the close approach happens inside a spherical region whose radius, called tidal radius, is given by: \\begin{equation} r_{td} = a_{B} \\left( \\frac {3 M_{P}}{m_{1}+m_{2}} \\right)^{1/3} \\label{eq:eqrtd} \\end{equation} where $a_{B}$, $M_{P}$, $m_{1}$ and $m_{2}$ are the binary semi-major axis, planet mass, primary and secondary asteroid masses, respectively. A possible outcome after disruption is the capture of one member of the primordial binary-asteroid. % The increasing number of binary-asteroid discoveries \\citep{noll06} \\gram{along with} the \\citet{agnorhamilton06} results, \\gram{have} motivated us to study the binary-asteroid capture process \\gram{in the context of} four-body dynamic\\gram{s}, where we considered Sun, Jupiter and a pair of asteroids. The purpose of this work is to identify the most important orbital characteristics inherent to the binary-asteroid capture/disruption process \\gram{that produce} the permanent capture of at least one member. \\gram{Compared to existing works }\\citep{agnorhamilton06, vokrouhlickyetal08}, \\gram{our study considers} the inclusion of solar perturbation. We found that the Sun's presence has a crucial influence on the binary capture/disruption process, at least in \\gram{the} planar case. This paper is built with the following structure: Section \\ref{sec:capturemodeling} describes the model we use in our study as well the adopted numerical approach. Section \\ref{sec:results} presents the results with an analysis of them. Finally, the last section summarizes our conclusions. ", "conclusions": "\\label{sec:conclusions} In this work we have studied the capture dynamics of binary-asteroids by looking for favorable conditions of capture. Our results present new perspectives about the problem of binary-asteroid captures emphasizing the importance of Sun's role in the dynamics. \\corr[31]{The Sun is not necessary to produce a binary rupture since Jupiter alone can do it. However, the Sun plays a key role in the disruption process in order to make the capture to become permanent.} The results allow us to comprehend about both binary-asteroid's features as well as intrinsic features of capture process' main stages. The observed characteristics at the first main stage, \\Ti, have revealed that: i) Tighter binary-asteroids are more susceptible to permanent captures than binaries with larger separation. In fact, the permanent capture probability behaves inversely proportional to the binary's semi-major axis. \\corr[NOVO]{This results indicate that binary's energy exchange allows the asteroid to become permanently captured. That is, since tighter binaries interactions are more intense, its members can exchange higher amount of energy. In such a way one asteroid can have its energy sufficiently decreased in order to not be able to escape from Jupiter.} \\corr[14]{Through this conclusion, we could argue that binary-asteroids with high eccentricities would disrupt more easily, but would not exchange the necessary amount of energy to result in a permanent capture.}\\corr[32]{As mentioned, there must exist a lower binary-separation limit below which the binary never disrupts and consequently there is no capture.} The observed characteristics at the second main stage, \\Tii, tell about process' features. It was shown that the angular position of bodies at disruption instant (\\Tii) are related with the permanent capture probability. Summarizing: ii) Disruption preferentially occurs when both asteroids are approximately aligned with Jupiter; Nevertheless, iii) disruptions which occurs when P2 is located between Jupiter and P1 result more often in permanent capture; Finally, we found that iv) the permanent capture probability is higher when the binary-asteroid disrupts in an \\ eastern quadrature\\, i.e., an angular position approximately 90\\dgr\\ after the binary cross the Sun-Jupiter direction. The permanent capture probability for an specific set of initial conditions, such derived from long time primary initial conditions, was shown to reach 20\\,\\%. Therefore, the good candidates are those derived from long capture time primary initial conditions. As a note, we give here a reference of a similar work submitted to Icarus journal, which also considers the solar perturbation. This work is available on astro-ph \\citep{philpottetal_arXiv09}. \\corr[33]{Finally, as the main goal of this paper was to address the favorable conditions which makes the permanent capture plausible, we have chosen a procedure to get initial conditions without taking into account where the incoming objects came from. It means that, in this work we have just tried the model plausibility without taking into account how it could reproduce the currently observed objects. So in order to get a more realistic probability, it must be considered several aspects as inclinations, mass ratios, binary eccentricities as well as to study how realistic are the trajectories of incoming objects.}" }, "1002/1002.2995_arXiv.txt": { "abstract": "{ A simple realization of inflation consists of adding the following operators to the Einstein-Hilbert action: $(\\partial\\phi)^2$, $\\lambda\\phi^4$, and $\\xi\\phi^2\\mathcal{R}$, with $\\xi$ a large non-minimal coupling. Recently there has been much discussion as to whether such theories make sense quantum mechanically and if the inflaton $\\phi$ can also be the Standard Model Higgs. In this work we answer these questions. Firstly, for a single scalar $\\phi$, we show that the quantum field theory is well behaved in the pure gravity and kinetic sectors, since the quantum generated corrections are small. However, the theory likely breaks down at $\\sim\\mpl/\\xi$ due to scattering provided by the self-interacting potential $\\lambda\\phi^4$. Secondly, we show that the theory changes for multiple scalars $\\vec\\phi$ with non-minimal coupling $\\xi\\vec\\phi\\cdot\\vec\\phi\\,\\mathcal{R}$, since this introduces qualitatively new interactions which manifestly generate large quantum corrections even in the gravity and kinetic sectors, spoiling the theory for energies $\\gtrsim \\mpl/\\xi$. Since the Higgs doublet of the Standard Model includes the Higgs boson and 3 Goldstone bosons, it falls into the latter category and therefore its validity is manifestly spoiled. We show that these conclusions hold in both the Jordan and Einstein frames and describe an intuitive analogy in the form of the pion Lagrangian. We also examine the recent claim that curvature-squared inflation models fail quantum mechanically. Our work appears to go beyond the recent discussions. } \\begin{document} ", "introduction": "\\label{sec:introduction} Cosmological inflation is our leading theory of the very early universe \\cite{Guth,Linde,AlbrechtSteinhardt82,Linde83}, although its underlying microphysics is still unknown. Slow-roll inflation occurs in many models constructed over the years, with a scattering of model-dependent predictions. Models of inflation are quite UV sensitive since it may have occurred at extremely high energy scales, far higher than that which we can probe at colliders, and since some models involve super-Planckian excursions in field space. This suggests that top-down approaches may be required to make progress, although progress in that direction has not been easy (e.g., see \\cite{Baumann:2007np}). On the other hand, it is interesting to explore simple models to see what we might learn and if we can connect inflation to low energy physics. One approach is to focus on dimension 4 Lagrangians, which allows the inclusion of the operators $(\\partial\\phi)^2$, $\\lambda\\phi^4$, and $\\xi\\phi^2\\mathcal{R}$ in addition to the Einstein Hilbert term $\\mpl^2\\mathcal{R}$ (by ``dimension 4\" we mean that each operator has a dimensionless coefficient, although gravity modifies the power counting since every term is an infinite tower of operators if we expand around flat space). The $\\xi\\phi^2\\mathcal{R}$ term is in fact required to exist for an interacting scalar field in curved space (although not for a Goldstone boson, for example). By taking the non-minimal coupling $\\xi$ to be large, a phase of inflation takes place, as originally discussed in \\cite{Salopek} with constraints discussed in \\cite{Fakir,Kaiser,Komatsu,Nozari:2007eq}. A large dimensionless coupling is unusual from the perspective of particle physics, leading to much recent debate as to whether such models make sense quantum mechanically \\cite{Salopek,Bezrukov,Barvinsky,Bezrukov2,Bellido,DeSimone:2008ei,Bez2008,Burgess,Espinosa,Bezrukov:2009db,Barvinsky:2009fy,Clark:2009dc,Lerner:2009xg,Barvinsky:2009ii,Figueroa:2009jw,Okada:2009wz,Einhorn:2009bh,Lerner:2009na,Mazumdar:2010sa}. In this work we show that for a single scalar field $\\phi$ the model is well behaved quantum mechanically in the pure gravity and kinetic sectors, since these do not generate large quantum corrections. However, the self-interacting potential does likely cause the theory to fail. For a vector of fields $\\vec\\phi$ carrying an $O(N)$ symmetry the theory manifestly generates large quantum corrections in the gravity and kinetic sectors and breaks down at high scales relevant to inflation. As far as we aware, this difference between the single field and multi-field case has not been fully appreciated in the literature. So multi-field models of non-minimal inflation are manifestly ruled out (including the Standard Model Higgs), while single field models of non-minimal inflation are also problematic, in the sense that the Einstein frame potential is non-polynomial and therefore these models likely fail due to high energy scattering processes. This presents a challenge to them having a UV completion. We also examine curvature-squared models and show that their scalar-tensor formulation is perturbatively well behaved in the gravity, kinetic, and potential sectors; this contradicts the claims of Ref.~\\cite{Burgess}. Our paper is organized as follows: in Section \\ref{sec:nonmin} we briefly review the class of non-minimal models, in Section \\ref{Quant} we discuss the quantum corrections, in Section \\ref{Pion} we discuss the pion analogy, in Section \\ref{Curvature} we discuss curvature-squared models, and finally we conclude in Section \\ref{Conclusion}. ", "conclusions": "Treated classically, a singlet scalar $\\phi$ with non-minimal coupling to gravity $\\xi\\phi^2\\mathcal{R}$ and $\\xi\\gg 1$ can drive a phase of slow-roll inflation. By examining the gravity and kinetic sectors alone, as in Ref.~\\cite{Burgess}, it is tempting to declare that such theories become strongly interacting at $\\mpl/\\xi$, which would spoil the model. Such a conclusion arises from power counting estimates of scattering processes in the Jordan frame without summing all diagrams to check for any possible cancellations. Here we have shown that cancellations do occur in the single field case, giving a Planckian cutoff. However, the potential sector is problematic: the Einstein frame potential is non-polynomial, varying on scales $\\Delta\\sigma\\sim\\mpl/\\xi$. Although its Coleman-Weinberg quantum corrections \\cite{Coleman} are small and the behavior of such a theory is considered an open problem in field theory, it is doubtful if high energy scattering amplitudes could be well behaved and if there is any sensible UV completion of the theory. This suggests that the true cutoff is $\\sim\\mpl/\\xi$ due to the potential sector. On the other hand, these power counting estimates in the gravity and kinetic sectors {\\em do} apply in the multi-field case: a vector of fields $\\vec\\phi$ is problematic as $\\phi_a+\\phi_b\\to\\phi_c+\\phi_d$ scattering becomes strong at energies $E\\gtrsim\\mpl/\\xi$ relevant to inflation (as noted in \\cite{Burgess,Atkins:2010eq}) and large quantum corrections are generated. This requires new physics to intervene or that something peculiar happens in the RG flow that rescues the theory, such as flow to a UV fixed point. But the latter scenario seems rather unlikely. Our conclusions were shown to be true in both the Jordan and Einstein frames. So if the Standard Model Higgs were a gauge singlet, then Higgs-inflation would be well behaved in the pure gravity and kinetic sectors, requiring a detailed analysis of the potential sector to draw conclusions about the quantum theory. However, since it is comprised of 4 real scalars in a complex doublet, it falls into the multi-field category and fails even at the level of tree-level scattering due to graviton exchange; breaking down due to Goldstone bosons. Explicitly identifying this difference between the single field case and the multi-field case was previously missed by us \\cite{DeSimone:2008ei} and others. Similarly, this may spoil other attempts to embed inflation into particle physics models with multiple scalar fields \\cite{Hertzberg:2007wc}, unless the underlying UV theory carries an appropriate structure. Our conclusions were drawn from an expansion around $\\langle\\phi\\rangle=0$. Although one might be concerned that this does not reflect the inflationary regime where $\\langle\\phi\\rangle$ is large, it does accurately reflect the end of inflation (the reheating era) in which $\\phi$ oscillates around $\\phi=0$. Hence our conclusions are quite robust in demonstrating that the reheating era cannot be described within effective field theory, i.e., that the effective Lagrangian must have large corrections as we {\\em approach} inflationary energy densities. Thus altering the inflationary Lagrangian itself. We also examined curvature-squared models $\\zeta\\,\\mathcal{R}^2$ \\cite{Starobinsky} with $\\zeta\\gg 1$ (although we consider such models to be ad hoc). They were recently criticized in Ref.~\\cite{Burgess} for failing to even make sense at scales $E\\gtrsim\\mpl/\\zeta^{1/3}$. Again there exists cancellation among tree-level diagrams, suggesting a Planckian cutoff. Furthermore, by formulating the quantum theory as a scalar-tensor theory in the Einstein frame (which makes the degrees of freedom manifest) we showed that the inflationary theory is well behaved. In this case, all scattering processes, whether they arise from the gravity or kinetic or potential sectors, satisfy unitarity bounds and generate very small quantum corrections for energies below $\\mpl$. The reason for this is that we have a singlet model with potential that varies on scales $\\Delta\\sigma\\sim\\mpl$ (as opposed to non-minimally coupled models which vary on scales $\\Delta\\sigma\\sim\\mpl/\\xi$.)" }, "1002/1002.1052_arXiv.txt": { "abstract": "We have used four telescopes at different longitudes to obtain near-continuous light curve coverage of the star HD~80606 as it was transited by its $\\sim$4-$M_{Jup}$ planet. The observations were performed during the predicted transit windows around the $25^{th}$ of October 2008 and the $14^{th}$ of February 2009. Our data set is unique in that it simultaneously constrains the duration of the transit and the planet's period. Our Markov-Chain Monte Carlo analysis of the light curves, combined with constraints from radial-velocity data, yields system parameters consistent with previously reported values. We find a planet-to-star radius ratio marginally smaller than previously reported, corresponding to a planet radius of $\\rp = 0.921 \\pm 0.036 R_{Jup}$. ", "introduction": "An important class of the $\\sim 400$ extra-solar planets known to date are the so-called ``hot Jupiters'', orbiting a fraction of an AU from their host star. Many of these have significantly larger radii than planets of similar mass in our Solar System. This has been a challenge for theoretical models of their structure. The intense radiation and tidal forces from the host star are expected to play a significant role \\citep[e.g.][]{Bodenheimer2001,Burrows2007}. For planets in highly eccentric orbits, both of these factors vary greatly over the course of the orbit. Thus, such planets provide interesting tests for models of the structure and dynamics of planetary atmospheres \\citep[e.g.][and references therein]{Irwin2008}. A planet in a 111-day orbit around the star HD~80606 was first detected in radial velocity observations \\citep{Naef2001}. The minimum mass (\\mpsini) of the planet is 3.9 times the mass of Jupiter. HD~80606~b has the most eccentric orbit of all the extra-solar planets known to date ($e = 0.93$). Infrared ($8 \\mu$m) observations clearly show the rapid heating of the planet's atmosphere during periastron passage \\cite[herein referred to as L09]{Laughlin2009}. Laughlin et al. also reported detection of a secondary eclipse, implying the orbital inclination is close to 90 degrees, and motivating efforts to observe a transit of the planet in front of the star. In early 2009, several groups observed a transit egress in photometry \\citep{Moutou2009,Garcia-Melendo2009,Fossey2009} and spectroscopy \\citep{Moutou2009}. Analyses of these data, together with old and new radial-velocity measurements provided the first constraints on the planet's radius and actual mass, as well as its orbital parameters (\\citealt{Pont2009}, herein referred to as P09; \\citealt{Gillon2009}, G09). These analyses are limited by the fact that the duration of the transit is not constrained, which in turn increases uncertainties in the system parameters. \\citet[W09]{Winn2009} report multiple observations of the transit ingress in June 2009, and combine their data with the earlier ingress observations and Keck radial-velocity data to constrain the transit duration. In this paper we present previously unpublished multi-site photometric observations of a complete transit of HD~80606~b on 14 February 2009, and an ingress on 25 October 2008. We describe the observations and the resulting data sets in section~\\ref{sec:obsphot}. In section~\\ref{sec:analysis} we describe the methods we used to analyse our data, and in section~\\ref{sec:results} we report and discuss the results. ", "conclusions": "\\label{sec:conclusion} We have obtained multi-site observations of a transit ingress and a complete transit of HD~80606~b across its host star. We analysed these data independently of any other photometric data, and found system parameters consistent with previously reported values. These observations were made using four telescopes at different sites. This allowed us to obtain near-continuous coverage of this 12-hour event. However, the differences between the instruments, telescopes and time-allocation procedures were, to an extent, limitations on our ability to obtain a uniform data set. In the near future, the completion of LCOGT's network of 1m robotic telescopes \\citep{Brown2010} will greatly facilitate observations of this kind, providing near-identical instrumentation at a number of sites, under the control of a flexible, central scheduling system." }, "1002/1002.4677_arXiv.txt": { "abstract": "\\ltarget\\ (hereafter \\target) is a compact white dwarf binary from the Sloan Digital Sky Survey that exhibits high-amplitude radial velocity variations on a period of $4.56\\,$hours. While an initial analysis suggested the presence of a neutron star or black-hole binary companion, a follow-up study concluded that the spectrum was better understood as a combination of two white dwarfs. Here we present optical spectroscopy and ultraviolet fluxes which directly reveal the presence of the second white dwarf in the system. \\target's spectrum is a composite, dominated by the narrow-lined spectrum from a cool, low gravity white dwarf ($T_{\\mathrm{eff}} \\simeq 6300\\,$K, $\\log g = 5$ to $6.6$) with broad wings from a hotter, high-mass white dwarf companion ($11,000$ to $14,000\\,$K; $\\sim 1\\,\\msun$). The high-mass white dwarf has unusual line profiles which lack the narrow central core to H$\\alpha$ that is usually seen in white dwarfs. This is consistent with rapid rotation with $v \\sin i = 500$ to $1750\\,\\kms$, although other broadening mechanisms such as magnetic fields, pulsations or a helium-rich atmosphere could also be contributory factors. The cool component is a puzzle since no evolutionary model matches its combination of low gravity and temperature. Within the constraints set by our data, \\target\\ could have a total mass greater than the Chandrasekhar limit and thus be a potential Type~Ia supernova progenitor. However, \\target's unusually low mass ratio $q \\approx 0.2$ suggests that it is more likely that it will evolve into an accreting double white dwarf (AM~CVn star). ", "introduction": "Close pairs of white dwarfs with combined masses greater than the Chandrasekhar limit have long been discussed as potential progenitors of Type~Ia supernovae \\citep{IbenTutukov1984,Webbink:DDs}. However, despite searches, \\citep{Robinson:DDs,Bragaglia:DDs,Marsh:friends,Napiwotzki:SPY} no secure examples of double white dwarfs both massive enough and short-period enough to merge within a Hubble time have been found. Only about 1 in a 1000 white dwarfs are required to be in such systems to match Type~Ia rates \\citep{Nelemans:DWDs}. Since only of order 1000 white dwarfs have been searched for binarity to date, the deficit is not yet significant, however it continues to be worth searching for more such systems. In an effort to do so we examined the spectra of DA white dwarfs (those showing spectra with hydrogen absorption only) from the SDSS survey \\citep{Eisenstein:wds}, looking for objects of discrepant radial velocity. One star, \\target\\ was an obvious outlier with a mean radial velocity of $-300\\,\\kms$. In this paper we present follow-up spectroscopy to elucidate the nature of this object. \\target\\ was the subject of a similar study by \\cite{Badenes:SDSS1257}. They found that it was a binary, measured a radial velocity semi-amplitude of $K_1 = 323 \\pm 6\\,\\kms$ on a period of $4.56\\,$hours and fitted their spectra with a white dwarf of temperature $\\sim 9000\\,$K and high mass, $\\sim 0.9\\,\\msun$. The period, radial velocity amplitude and white dwarf mass give a minimum mass for the companion of $1.6\\,\\msun$, suggesting that it is either a neutron star or black-hole. \\cite{Badenes:SDSS1257} estimated a distance of $48\\,$pc for \\target, implying that such systems may be rather common. \\begin{figure*} \\centering \\hspace*{\\fill} \\includegraphics[width=0.9\\textwidth]{f1.eps} \\hspace*{\\fill} \\caption{The panels show phase-folded trailed spectra of the first four Balmer lines of \\target, H$\\alpha$--$\\delta$, left to right. The left-hand panel shows the raw data; the right-hand panel, which shows the data after removal of the primary's motion and mean spectrum, reveals anti-phased features from the massive secondary white dwarf. (Note the change of horizontal scale between the two panels.) \\label{fig:pbin_trail}} \\end{figure*} \\cite{Badenes:SDSS1257} recognised that their spectral fits were problematic, and in particular failed to fit the narrow cores of the Balmer lines. Thus while they favored a neutron star or a black-hole for the companion, they could not entirely eliminate the possibility that it was another white dwarf. \\cite{kulkarni+vankerkwijk10-1} took three high signal-to-noise spectra which showed asymmetries suggesting exactly this; they showed that their spectra could be fit by a combination of a cool, low mass white dwarf ($6250 \\pm 250\\,$K, $0.15\\pm 0.05\\,\\msun$) plus a hotter, massive white dwarf companion ($13$,$000\\pm 800\\,$K, $0.92 \\pm 0.13\\,\\msun$). \\cite{kulkarni+vankerkwijk10-1}'s paper appeared shortly after the original submission of our work. Although both our papers agreed upon the basic double white dwarf nature of \\target, there were significant differences of detail, with inconsistencies in both masses and temperatures. In an effort to understand these, we have since carried out additional fits to our data, which we present here along with our original approach. In addition, an improved \\emph{Swift} calibration has given us more confidence in the modelling of the UV-optical spectral energy distribution. Our new results agree more closely with \\cite{kulkarni+vankerkwijk10-1} than our original analysis, but also reveal uncertainties in the system properties which make it impossible to establish securely such fundamental properties as whether the system is super-Chandrasekhar or not. We begin with a description of our observational material. ", "conclusions": "We find that the putative white dwarf--black-hole/neutron star binary, \\target, is a double white dwarf. \\target\\ is composed of a very low mass $\\sim 0.2\\,\\msun$ white dwarf together with an extremely massive ($> 1\\,\\msun$) white dwarf. As long as the massive white dwarf avoids accretion-induced collapse or explosion, \\target\\ will evolve into a hydrogen-deficient accreting binary star, but may later explode as a sub-luminous Type~Ia. The massive white dwarf shows signs of rapid rotation ($v \\sin i > 500 \\,\\kms$) which suggests that the most recent phase of mass transfer might not have involved a common envelope, contrary to current models of double white dwarf populations. Some inconsistencies in the parameters of the two white dwarfs remain that are probably caused by difficulties in fitting their blended spectra; further observations are required to clarify these." }, "1002/1002.3327_arXiv.txt": { "abstract": "We develop a method for calculating the equilibrium properties of the liquid-solid phase transition in a classical, ideal, multi-component plasma. Our method is a semi-analytic calculation that relies on extending the accurate fitting formulae available for the one-, two-, and three-component plasmas to the case of a plasma with an arbitrary number of components. We compare our results to those of Horowitz, Berry, \\& Brown (Phys.~Rev.~E {\\bf 75}, 066101, 2007), who use a molecular dynamics simulation to study the chemical properties of a 17-species mixture relevant to the ocean-crust boundary of an accreting neutron star, at the point where half the mixture has solidified. Given the same initial composition as Horowitz et al., we are able to reproduce to good accuracy both the liquid and solid compositions at the half-freezing point; we find abundances for most species within $10\\%$ of the simulation values. Our method allows the phase diagram of complex mixtures to be explored more thoroughly than possible with numerical simulations. We briefly discuss the implications for the nature of the liquid-solid boundary in accreting neutron stars. ", "introduction": "\\label{sec:intro} During the crystallization of a plasma containing multiple ion species, the chemical composition of the solid is in general different from that of the liquid. This type of chemical separation is important for both white dwarfs \\citep{hansen03} and accreting neutron stars \\citep{horowitz07}. The interior of a white dwarf is a mixture of carbon, oxygen, and traces of other elements, most abundantly neon. As the star cools, chemical separation leads to the formation of an oxygen- and neon-rich core. The energy released through the gravitational settling of the denser core material heats the star and can delay cooling by several Gyr \\citep{isern91}. A neutron star accretes mostly hydrogen and helium from its companion, but this material undergoes a series of nuclear reactions, including rapid proton capture \\citep{schatz01} and then electron capture reactions \\citep{gupta07}, to produce a variety of elements. Through accretion the mixture is pushed deep into the star and solidifies. Recent numerical simulations have shown that the mixture undergoes chemical separation during solidification \\citep{horowitz07}, possibly forming a two-phase solid \\citep{horowitz09b}. The composition of the liquid ocean and the structure and composition of the crust have important implications for a range of observed phenomena. For example, the resulting thermal conductivity determines the cooling rate of transiently accreting neutron stars following extended accretion outbursts \\citep{shternin07,brown09}. The mechanical strength of the crust limits the size of a possible crust quadrupole and therefore gravitational wave emission \\citep{horowitzkadau09}. Several groups have studied the liquid-solid phase transition and chemical separation of two- and three-component plasmas in the classical, ideal limit (i.e., ignoring quantum mechanical effects on the ions and treating the electrons as a uniform background; cf. Ref.~\\cite{potekhin00}). Early works (e.g., Ref.~\\cite{mochkovitch83}) studied phase transitions in carbon-oxygen plasmas, but the approximations used were too crude for application to the interiors of white dwarfs. Accurate calculations using the mean spherical approximation in the density-functional formalism were performed by \\citet{barrat88}, who studied carbon-oxygen plasmas, and by \\citet{segretain93}, who studied arbitrary two-component plasmas with atomic number $Z$ ratios up to 2 (see also Ref.~\\cite{segretain96}, where carbon-oxygen-neon plasmas are examined). Using Monte Carlo calculations and $Z$ ratios up to 5, \\citet{ogata93} studied arbitrary two- and three-component plasmas and \\citet{dewitt03} studied arbitrary two-component plasmas with a very accurate measurement of the liquid free energy (see also Refs.~\\cite{iyetomi89,dewitt96}). All of these groups present phase diagrams as a function of ion abundance, and some \\citep{ogata93,dewitt03} also present fitting formulae for the liquid and solid free energies. Using these diagrams and fitting formulae, one can determine the phase transition properties for a two-component plasma of any ion type and abundance. These calculations are particularly useful for the interior of a white dwarf, where there are only two or three dominant elements. But in the ocean of an accreting neutron star there are around 10-20 elements with abundances $> 1\\%$ \\citep{gupta07}, each one with a potentially important effect on the behavior of the phase transition and chemical separation of the mixture. The available analytic or numerical results for this type of system are extremely limited. We are aware of only one study of phase transitions in plasmas with more than three components, that of \\citet{horowitz07} (see also Refs.~\\cite{horowitz09a,horowitz09b}). These authors used molecular dynamics simulations to study a 17-component plasma with a composition similar to that expected at the ocean-crust interface of an accreting neutron star. Due to the large amount of computing power necessary to run each simulation, the phase transition properties have so far only been calculated for one composition. We present here a method for rapidly calculating the properties of the liquid-solid phase transition in a multi-component plasma in the classical ideal limit, for any initial composition and ion types. Our method is a semi-analytic calculation that relies on extending the accurate fitting formulae available for the one-, two-, and three-component plasmas to the case of a plasma with an arbitrary number of components. We test our method using the one data point available for a plasma with more than three components, the calculation of \\citet{horowitz07}, and show that it performs very well in that specific case. The paper is organized as follows. In Section~\\ref{sec:method} we describe the semi-analytic calculation as it applies to the one-component plasma (Section~\\ref{sub:ocp}), the two-component plasma (Section~\\ref{sub:tcp}), and the multi-component plasma (Section~\\ref{sub:mcp}). In Section~\\ref{sec:results} we present our results for the 17-component mixture of \\citet{horowitz07}. We conclude in Section~\\ref{sec:discuss}. The pressure term in the Gibbs free energy and its effect on the phase transition, the importance of the deviation from linear mixing for the liquid free energy, and a simplified derivation of the deviation from linear mixing for the solid free energy, are discussed in three appendices. ", "conclusions": "\\label{sec:discuss} Using results from simulations of one-, two-, and three-component plasmas, we have developed a method for calculating the equilibrium properties of the liquid-solid phase transition in a plasma with an arbitrary number of components, in the approximation of a classical ion plasma in a uniform electron background. We used this method to calculate the phase transition properties for a 17-component plasma with a composition similar to that which might exist in the ocean of an accreting neutron star, and compared the results to those of a molecular dynamics simulation done at the same composition (HBB \\cite{horowitz07}). We found that our method accurately reproduces the results of the HBB simulation. Two sources of error in the simulation may mean that our results represent the actual system even more accurately than this comparison suggests: First, the finite size of the simulation introduces statistical errors which for some components are larger than the discrepancies between the two works. Second, the system is still evolving at the end of the simulation, with many components approaching the values predicted by our calculation. As in the simulation of HBB, we have followed the 17-component mixture until it reaches the state of $50\\%$ liquid and $50\\%$ solid. Under these conditions, the term representing the deviation from the linear mixing rule for the solid, $\\Delta f_s$, is a perturbation on the other terms in the free energy of the solid [see Eq.~(\\ref{eq:fsMCP})]. In principle our calculation can continue to larger fractions of solid, i.e., larger values of the Coulomb coupling parameter $\\Gamma$. However, because $\\Delta f_s$ increases linearly with $\\Gamma$ and eventually dominates the free energy, the calculation at $\\Gamma$ above the half-freezing point is more sensitive to the form chosen for $\\Delta f_s$. There is some numerical confirmation of our simple approximation for $\\Delta f_s$, Eqs.~(\\ref{eq:delvs}) and (\\ref{eq:delusMCP}), for two- and three-component mixtures at large $\\Gamma$, but only for a very limited set of parameters (see Ref.~\\cite{ogata93}). Further numerical simulations are necessary to test the validity of these equations at large $\\Gamma$ for general parameters and $(m>3)$-component plasmas. Another consequence of the large and positive $\\Delta f_s$ term is that for certain compositions, it is energetically favorable for a single solid phase to separate into two or more solid phases (see Section~\\ref{sub:tcp}). Such a phase separation occurs at large $\\Gamma$ in the 17-component plasma simulated by \\citet{horowitz09a}. With our calculation we have not yet found any two-solid mixtures that represent the lowest energy state of the HBB plasma, in part because the shape of free energy surface for the solid phase is very complicated at large $\\Gamma$. We leave a more careful study of the solid-solid unstable region for future work. Once these issues are resolved, our calculation will allow the complete phase diagram of multi-component mixtures to be determined. We expect that these results will have important implications for the structure of the liquid-solid boundary in accreting neutron stars. For example, for an ocean temperature of $T=10^8 T_8~{\\rm K}$, an O-Se mixture with the same proportion of oxygen and selenium as in the HBB mixture (i.e., $\\sim 10\\%$-$90\\%$) will begin to freeze at a density of $\\rho \\simeq 2\\times 10^7 T_8^3 (\\mu_e/2)~{\\rm g/cm^3}$, where $\\mu_e$ is the mean molecular weight per electron. Assuming that accretion is slow enough that the liquid and solid can come into equilibrium at each depth, our phase diagram for a charge ratio $R_Z=34/8$ in Fig.~\\ref{fig:gibbscomp} (or Fig.~\\ref{fig:delflcomp}) shows that the mixture will reach $50\\%$ solid within a factor of two in density, but that complete freezing will not occur until much deeper, by a factor of $\\simeq (34/8)^5 \\simeq 1400$ in density (corresponding to $\\rho \\simeq 3\\times 10^{10} T_8^3~{\\rm g/cm^3}$). This is a very different picture than the sharp transition between liquid and solid expected for a one-component plasma, and assumed in previous work on accreting neutron stars. Further work is needed to understand the effects of the various time-dependent processes that are active concurrent with accretion in the ocean-crust transition layer, such as crystallization, diffusion, and sedimentation. For example, sedimentation of the heavier solid particles could be important at low accretion rates, narrowing the transition layer." }, "1002/1002.0691_arXiv.txt": { "abstract": "It has been known for nearly three decades that high redshift radio galaxies exhibit steep radio spectra, and hence ultra-steep spectrum radio sources provide candidates for high-redshift radio galaxies. Nearly all radio galaxies with z $>$ 3 have been found using this redshift-spectral index correlation. We have started a programme with the Giant Metrewave Radio Telescope (GMRT) to exploit this correlation at flux density levels about 10 to 100 times deeper than the known high-redshift radio galaxies which were identified primarily using the already available radio catalogues. In our programme, we have obtained deep, high resolution radio observations at 150 MHz with GMRT for several $'${\\it deep}$'$ fields which are well studied at higher radio frequencies and in other bands of the electromagnetic spectrum, with an aim to detect candidate high redshift radio galaxies. In this paper we present results from the deep 150 MHz observations of LBDS-Lynx field, which has been already imaged at 327, 610 and 1412 MHz with the Westerbork Synthesis Radio Telescope (WSRT) and at 1400 and 4860 MHz with the Very Large Array (VLA). The 150 MHz image made with GMRT has a rms noise of $\\sim$ 0.7 mJy/beam and a resolution of $\\sim 19^{''} \\times 15^{''}$. It is the deepest low frequency image of the LBDS-Lynx field. The source catalog of this field at 150 MHz has about 765 sources down to $\\sim$ 20\\% of the primary beam response, covering an area of about 15 degree$^2$. Spectral index was estimated by cross correlating each source detected at 150 MHz with the available observations at 327, 610, 1400 and 4860 MHz and also using available radio surveys such as WENSS at 327 MHz and NVSS and FIRST at 1400 MHz. We find about 150 radio sources with spectra steeper than 1. About two-third of these are not detected in Sloan Digital Sky Survey, hence are strong candidate high-redshift radio galaxies, which need to be further explored with deep infra-red imaging and spectroscopy to estimate the redshift. ", "introduction": "Since powerful radio sources reside in massive ellipticals (Best et al. 1998), finding radio galaxies at high redshifts is important to understand the radio evolution of galaxies. It had been noticed in the early 80's that the fraction of radio sources that can be optically identified is lower by a factor of 3 or more for steep spectrum radio sources ($\\alpha > 1; S_\\nu \\propto \\nu^{-\\alpha}$; Tielens, Miley \\& Wills, 1979, Blumenthal \\& Miley, 1979; Gopal-Krishna \\& Steppe, 1981), suggesting that the radio sources with steeper spectra at decimeter wavelengths are more distant compared to the ones with normal spectra ($ 1 > \\alpha > 0.5$). Over the years, this correlation has been exploited to search for radio sources at high redshifts (Rottgering et al., 1994; De Breuck et al. 2000, 2002; Klamer et al. 2006). Since the radio emission does not suffer from dust absorption, selecting candidate high redshift radio galaxies (HzRGs) at radio frequencies provides an optically unbiased sample. The most distant radio galaxy at z = 5.19 was discovered using the spectral index - redshift correlation (van Breugel et al. 1999). Until today, about 45 radio galaxies are known beyond redshift of 3, and nearly all of them were discovered through the radio spectral index - redshift correlation. Thus, this is the most efficient method to find HzRGs. In addition to constraining models of formation and evolution of massive galaxies, HzRGs can also be used to study the environment at those epochs. Deep radio polarimetric observations of several HzRGs have shown large rotation measures (RM) of the order of thousands of rad m$^{-2}$ (Pentericci et al. 2000; Carilli et al. 1997; Athreya et al. 1998; Kronberg et al. 2008). Pentericci et al. (2000) also found that the fraction of radio galaxies with extreme Faraday rotation is increasing with redshift which is consistent with the hypothesis that the environment is denser at high redshift. Since at moderate redshifts large RMs are known to occur in radio galaxies residing near the center of clusters of galaxies, the HzRGs are an excellent tool to study the (proto) cluster environment at high redshifts. It is therefore important to identify and study as many HzRGs as possible. A major programme to search for HzRGs using the radio spectral index - redshift correlation was carried out in the last decade using different radio surveys at 1400 MHz and lower frequencies (eg: Rottgering et al. 1994, De Breuck et al. 2000, 2002, 2004, 2006; Klamer et al. 2006; see also review by Miley and De Breuck, 2008). The searches were mostly limited by the sensitivity limit of these shallow, wide-area radio surveys. This bias has led to the detection of only the brightest HzRGs which are at the top end of the radio luminosity function (See section 2), which is about three orders of magnitude brighter than the luminosities at the lower end of FR-II population (Fanaroff \\& Riley, 1974). Since steep spectrum radio sources are preferentially selected at low radio frequencies, deeper observations at these frequencies are needed to discover HzRGs that are 10 to 100 times less luminous than most of the known HzRGs, and these could be the high-redshift counterparts of more typical FR-II radio galaxies. The 150 MHz band of GMRT (Giant Metrewave Radio Telescope, India, http://www.ncra.tifr.res.in) with its large field of view (half-power beam width of 3 degrees), high angular resolution ($\\sim 20^{''}$) and better sensitivity ($\\sim$ 1 mJy from a full synthesis observation) is well suited for this kind of work. For example, a steep spectrum radio source ($\\alpha > 1$) at the completeness limit of WENSS at 327 MHz (Rengelink et al. 1997), will have a flux density $>$ 65 mJy at 150 MHz. Our experience with GMRT at 150 MHz shows that it is possible to reliably detect sources more than 10 times fainter than this value. The large field of view of GMRT also enables us to detect close to thousand radio sources in a single pointing. Encouraged by this, we have started a programme to observe several carefully chosen deep fields at 150 MHz with GMRT with an aim to detect steep spectrum radio sources to flux density levels much fainter than that of known high-redshift radio sources (see Section 2). As the sample at low radio frequencies goes deeper, the available surveys at higher radio frequencies such as FIRST won't be deep enough to get the spectral index estimates, particularly for steep spectrum sources. For example, from the present work, it is seen that if we use only the FIRST catalogue, 37\\% of the sources below 18 mJy at 150 MHz (which corresponds to 1 mJy at 1.4 GHz for a spectral index of 1.3) do not have counter parts in FIRST as against 9\\% for stronger sources. This limitation can be addressed by deeper surveys of this region at higher radio frequencies. Therefore, we have chosen the well known deep fields because significant amount of data already exist for them at higher radio frequencies and in the optical and infrared bands. These deep observations will allow us to estimate the spectral index for faint sources which are below the detection limit of NVSS or FIRST. The deep optical data will help to estimate the redshifts and eliminate nearby and known objects among the steep spectrum sources. Wherever needed, additional deep observations with the GMRT at 610 MHz will be obtained for a better estimate of radio spectra. In addition to the steep spectrum radio sources, this kind of survey will also help us to understand the evolution (LogN-LogS) and spectral index properties of faint radio sources, to detect low-power FRI sources, relic emission and large angular size radio sources. In this paper we present deep 150 MHz radio observations of the LBDS-Lynx field with the GMRT with the primary aim of detecting steep spectrum radio sources which are candidate HzRGs. The Leiden-Berkeley Deep Survey (LBDS) in the Lynx area (Windhorst et al. 1984) was carried out in the mid eighties to better understand the nature of faint radio sources and their cosmological evolution. Multi-band optical data were obtained with the Kitt Peak Mayall 4m telescope and subsequently deep radio observations were carried out at 327 MHz and 1412 MHz with WSRT, complimented by 1400 MHz and 4860 MHz imaging with the VLA (Windhorst et al 1984, 1985, Oort, 1987, Oort et al. 1988). These radio observations showed a moderate increase in the number of radio sources below $\\sim$ 5 mJy at 1.4 GHz, which were predominantly identified with blue galaxies, a population responsible for the upturn in the source counts at faint flux density levels (Windhorst et al. 1985). The spectral index studies of radio sources selected at 327 MHz in this region using the 327 MHz and 1400 MHz data showed that the median spectral index was flatter for fainter sources selected at 327 MHz (Oort et al. 1988). The sources with the steepest spectra ($\\alpha > 1$) were mostly unidentified in the optical 4m telescope images, once again indicating that they were likely to be very distant. The sources we detected at 150 MHz with the GMRT were compared with the deep observations of this region at 327, 610 and 1412 MHz with the WSRT and at 4860 MHz with the VLA. We have also used the available catalogues such as WENSS at 327 MHz and NVSS and FIRST at 1400 MHz to estimate spectral indices. We identify about 150 steep spectrum radio sources from this field. About two-third of them are not detected to the limit of SDSS, hence are strong candidate HzRGs. In Section 2, we present supporting arguments for reviving the search for HzRGs using the steep spectrum technique. Observations, data analysis and the data obtained are presented in Section 3. Results and discussions are presented in Section 4 and concluding remarks in Section 5. ", "conclusions": "The radio spectral-index redshift correlation is the most efficient method used to detect high-redshift radio galaxies. Up to now, about 45 radio galaxies beyond redshift of 3 have been discovered using this method. We have shown that this known high-z population is just the tip of the iceberg, in the sense that they are the most luminous objects in this class. Radio sources which are 10 to 100 times less luminous than these are yet to be discovered, and these are essential to fully understand the cosmological evolution of radio galaxies. We have initiated a major programme to search for steep spectrum radio sources with the GMRT with the aim to detect high-redshift radio galaxies of moderate luminosity. The fields for this programme are carefully chosen such that extensive data already exist at higher radio frequencies and deep optical imaging and/or spectroscopy is also available for most of the fields. Here we have presented the results from deep 150 MHz low frequency radio observations with the Giant Metrewave Radio Telescope (GMRT), India reaching an rms noise of of $\\sim$ 0.7 mJy/beam and a resolution of $\\sim$ 20$^{''}$. Further, several deep observations exist, mainly at 327, 610 and 1412 MHz for this field. The radio spectral index analysis of the sources in the field was done using these deep observations and also using the WENSS at 325 MHz and the NVSS and FIRST at 1400 MHz. We have demonstrated that this GMRT survey can search for high-redshift radio galaxies more than an order of magnitude fainter in luminosity compared to most of the known HzRGs. We provide a sample of about 100 candidate HzRGs with spectral index steeper than 1 and no optical counterpart to the SDSS sensitivity limits. A significant fraction of the sources are compact, and are strong candidates for high-redshift radio galaxies. These sources will need to be followed up at optical and near-infrared bands to estimate their redshifts." }, "1002/1002.0833_arXiv.txt": { "abstract": "{We provide a derivation from first principles of the primordial bispectrum of scalar perturbations produced during inflation driven by a canonically normalized scalar field whose potential exhibits small sinusoidal modulations. A potential of this type has been derived in a class of string theory models of inflation based on axion monodromy. We use this model as a concrete example, but we present our derivations and results for a general slow-roll potential with superimposed modulations. We show analytically that a resonance between the oscillations of the background and the oscillations of the fluctuations is responsible for the production of an observably large non-Gaussian signal. We provide an explicit expression for the shape of this \\textit{resonant non-Gaussianity}. We show that there is essentially no overlap between this shape and the local, equilateral, and orthogonal shapes, and we stress that resonant non-Gaussianity is \\textit{not} captured by the simplest version of the effective field theory of inflation. We hope our analytic expression will be useful to further observationally constrain this class of models.} ", "introduction": "The study of anisotropies in the cosmic microwave background radiation over the past two decades has dramatically improved our understanding of the early universe. There is now strong evidence that the anisotropies we see today originated from primordial fluctuations generated in the very early universe, and we have learned that these primordial fluctuations have a nearly scale-invariant spectrum. Furthermore, the data still contains no evidence for a deviation from adiabaticity, or Gaussianity~\\cite{Komatsu:2010fb}. Even though the case is by no means closed, these properties of the primordial fluctuations certainly support the idea that they originated as quantum fluctuations during inflation, a phase of nearly exponential expansion of the universe~\\cite{Guth:1980zm}. While observations are now good enough to rule out some of the simplest inflationary models involving only a single slowly rolling field with canonical kinetic term, other models in this class are still compatible with all existing data. These models predict an adiabatic spectrum with primordial non-Gaussianities that are too small to be observed, but this is not a generic prediction of inflation. Many models of inflation have been constructed and studied that can lead to an observable departure from Gaussianity. Some popular possibilities are models with multiple fields, non-canonical kinetic terms, light spectator fields and a violation of slow-roll. Beyond an existence proof that observably large non-Gaussianities can be generated, these models provide us with useful theoretical expectations to guide our search. For a Gaussian signal, all odd $n$-point functions vanish and the higher even $n$-point functions are given in terms of sums of products of the two-point function. The most straightforward way to look for a departure from Gaussianity is then to look for a non-zero three-point function. In Fourier space the three-point function depends on three momenta. Translational invariance of the background geometry ensures that these momenta add up to zero and thus form a triangle. Rotational invariance furthermore dictates that the three-point function can only depend on the three independent scalar products of these momenta. The information contained in the three-point function can thus be captured by a function of three variables, that can be thought of as two angles and one side of the triangle. Since the dependence is a priori completely arbitrary, a model independent measurement would be desirable and would provide a precious criterion to discriminate between otherwise indistinguishable models. Unfortunately, progress in this direction is very hard. (For a review see {\\it e.g.}~\\cite{Liguori:2010hx}). Essentially all phenomenological analyses start from some explicit form of the three-point function guided both by theoretical expectations and by the simplicity of the numerical analysis necessary to compare it with the data. (See, however,~\\cite{Fergusson:2009nv}.) Once a ``shape'' has been chosen, only the amplitude of this type of non-Gaussianity remains as a parameter, which is conventionally called $f^\\text{shape}$. So far only a handful of scale-invariant shapes have been looked for in the data. The most recent observational bounds on the magnitude for various shapes from the 7-year WMAP data at 95\\% CL are~\\cite{Komatsu:2010fb}: \\begin{center} \\begin{tabular}{lc} local non-Gaussianity & $-100$. However, models where the synchrotron peak occurs at frequencies $\\nu<\\nu_{NIR}$ are not within the 90\\% parameter confidence region when fitting our numerical models to the NIR and X-ray data (we fit the models and derived confidence regions using the X-ray spectral fitting package XSPEC). Furthermore, blue indices ($\\beta_{NIR}>0$) for intermediate to bright emission from Sgr A* are favoured not just by our observational constraint on the MIR-NIR spectral index, but by evidence from other observations as well \\citep{gen03,gil06,hor07}. There is much still to learn from the spectacular multiwavelength observation of April 4, 2007. The simultaneous lightcurves show intriguing differences: the L'-band lightcurve is broader (has longer duration) than the X-ray lightcurve, and it also has pronounced substructures while the X-ray lightcurve is comparatively smooth. These properties are not straightforward to understand in the context of any emission model. In the SSC scenario a shorter duration X-ray flare can arise naturally without changes in the magnetic field, due to the quadratic dependence of the SSC luminosity on the synchrotron luminosity, but there is no obvious solution to the substructure problem. For the cooling break synchrotron and the submm IC models fluctuations in magnetic field together with a general decrease in magnetic field during the flare could explain the substructure and duration problems, since in both cases the NIR emission is dependent on the magnetic field while the X-ray emission is not. Taking the spectral properties also into account, so far the cooling break synchrotron model appears to be the most promising model to explain the April 4, 2007 flare. Time dependent modeling of the lightcurves will give us further insights into the flare mechanism and the dynamics of their production. \\vspace{-0.15cm}" }, "1002/1002.2627_arXiv.txt": { "abstract": "{ The main goal of the Pennsylvania - Toru\\'n Planet Search (PTPS) is detection and characterization of planets around evolved stars using the high-accuracy radial velocity (RV) technique. The project is performed with the 9.2 m Hobby-Eberly Telescope. To determine stellar parameters and evolutionary status for targets observed within the survey complete spectral analysis of all objects is required. In this paper we present the atmospheric parameters (effective temperatures, surface gravities, microturbulent velocities and metallicities) of a subsample of Red Giant Clump stars using strictly spectroscopic methods based on analysis of equivalent widths of Fe~I and Fe~II lines. It is shown that our spectroscopic approach brings reliable and consistent results. } ", "introduction": "\\label{RG-intro} Over 350 extrasolar planets are known today and most of them were discovered by the radial velocity (RV) technique. Apart from many surveys searching for exoplanets around solar-like dwarfs, a few surveys looking for planetary companions orbiting evolved stars exist. For proper interpretation of the results obtained from RV studies of late-type giants detailed knowledge of their physical parameters is required (e.g. \\cite{S06}, \\cite{H07}, \\cite{T08}). Determination of the atmospheric parameters with high accuracy allows us to place each star on the Hertzsprung-Russell diagram (HRD) and better understand the formation and evolution of the object and its companion. In Fig.~\\ref{RG-fig1} we present 201 objects from the Pennsylvania - Toru\\'n Planet Search (PTPS), red giants from the Red Giant Clump (RGC) and approximately 11~000 stars from the Hipparcos and Tycho catalogues (gray dots). The effective temperatures and luminosities were taken from \\cite{A08}. \\begin{figure} \\center \\includegraphics[width=7cm]{HRD-201.eps} \\caption{HRD for 201 red giants from our sample. Red circles indicate five planetary system hosts from PTPS (\\cite{N07}, \\cite{N09a}, \\cite{N09b}).} \\label{RG-fig1} \\end{figure} ", "conclusions": "" }, "1002/1002.2411_arXiv.txt": { "abstract": "% The availability of continuous helioseismic data for two consecutive solar minima has provided a unique opportunity to study the changes in the solar interior that might have led to this unusual minimum. We present preliminary analysis of intermediate-degree mode frequencies in the 3 mHz band during the current period of minimal solar activity and show that the mode frequencies are significantly lower than those during the previous activity minimum. Our analysis do not show any signature of the beginning of cycle 24 till the end of 2008. In addition, the zonal and meridional flow patterns inferred from inverting frequencies also hint for a delayed onset of a new cycle. The estimates of travel time are higher than the previous minimum confirming a relatively weak solar activity during the current minimum. ", "introduction": "The delayed onset of solar cycle 24 and the prolonged period of minimal solar activity have invoked lots of interest in a variety of studies that might be useful to characterize the sun in a quiet state. Studies based on the helioseismic data have provided conflicting estimates of the length of previous cycle and shown that the present minimum is indeed the deepest in many aspects \\citep{bison09, howe09, david09,sct09b}. Since acoustic modes spend most of the time in the outer layers of the solar interior, the intermediate- and high-degree modes can be useful in interpreting the conditions in the convection zone. In this context, we investigate the response of these modes to the period of minimum activity and compare the response with the previous one. ", "conclusions": "" }, "1002/1002.3598_arXiv.txt": { "abstract": "We investigate the clustering properties of $\\sim$1550 broad-line active galactic nuclei (AGNs) at $\\langle$$z$$\\rangle$=0.25 detected in the {\\em ROSAT} All-Sky Survey (RASS) through their measured cross-correlation function with $\\sim$46000 Luminous Red Galaxies (LRGs) in the Sloan Digital Sky Survey. By measuring the cross-correlation of our AGN sample with a larger tracer set of LRGs, we both minimize shot noise errors due to the relatively small AGN sample size and avoid systematic errors due to the spatially-varying Galactic absorption that would affect direct measurements of the auto-correlation function (ACF) of the AGN sample. The measured ACF correlation length for the total RASS-AGN sample ($=1.5 \\times 10^{44}$ erg\\,s$^{-1}$) is $r_{\\rm 0}=4.3^{+0.4}_{-0.5}$ $h^{-1}$ Mpc and the slope $\\gamma=1.7^{+0.1}_{-0.1}$. Splitting the sample into low and high $L_{\\rm X}$ samples at $L_{\\rm X}^{0.5-10\\,{\\rm keV}} = 10^{44}$ erg\\,s$^{-1}$, we detect an X-ray luminosity-dependence of the clustering amplitude at the $\\sim$2.5$\\sigma$ level. The low $L_{\\rm X}$ sample has $r_{\\rm 0}=3.3^{+0.6}_{-0.8}$ $h^{-1}$ Mpc ($\\gamma=1.7^{+0.4}_{-0.3}$), which is similar to the correlation length of blue star-forming galaxies at low redshift. The high $L_{\\rm X}$ sample has $r_{\\rm 0}=5.4^{+0.7}_{-1.0}$ $h^{-1}$ Mpc ($\\gamma=1.9^{+0.2}_{-0.2}$), which is consistent with the clustering of red galaxies. From the observed clustering amplitude, we infer that the typical dark matter halo (DMH) mass harboring RASS-AGNs with broad optical emission lines is log $(M_{\\rm DMH}/(h^{-1}\\,M_\\odot)) =12.6^{+0.2}_{-0.3}$, 11.8$^{+0.6}_{-\\infty}$, 13.1$^{+0.2}_{-0.4}$ for the total, low $L_{\\rm X}$, and high $L_{\\rm X}$ RASS-AGN samples, respectively. ", "introduction": "Galaxies and active galactic nuclei (AGNs) are not distributed randomly in the universe. The small primordial fluctuations in the matter density field present in the very early universe have progressively grown through gravitational collapse to create the complex network of clusters, groups, filaments, and voids seen in the distribution of structure today. Galaxies and AGNs, as well as groups and clusters of galaxies, are believed to populate the collapsed dark matter halos (DMHs). The clustering of galaxies and AGNs therefore reflects the spatial distribution of dark matter in the universe. This allows clustering measurements to be used to derive cosmological parameters (e.g., \\citealt{peacock_cole_2001}; \\citealt{abazajian_zheng_2005}). However, these measurements also allow us to study the complex physics which governs the creation and evolution of galaxies and AGNs, as well as the co-evolution of galaxies and AGNs. The co-evolution scenario is motivated by the observed correlation between the mass of the central super-massive black hole (SMBH) and the stellar velocity dispersion in the bulge (\\citealt{gebhardt_bender_2000}; \\citealt{ferrarese_merritt_2000}), lending strong evidence to an interaction or feedback mechanism between the SMBH and the host galaxy. The specific form of the feedback mechanism, as well as the details of the AGN triggering, accretion, and fueling mechanisms, remains unclear. Different cosmological simulations address possible scenarios for the co-evolution of AGNs and their host galaxies (e.g., \\citealt{kauffmann_haehnelt_2000}; \\citealt{dimatteo_springel_2005}; \\citealt{cattaneo_dekel_2006}). Large volume, high-resolution simulations that include physical prescriptions for galaxy evolution and AGN feedback make predictions for the spatial clustering and large-scale environments of AGNs and galaxies (\\citealt{springel_white_2005}; \\citealt{colberg_dimatteo_2008}; \\citealt{bonoli_marulli_2009}). Observed clustering measurements of AGNs can be used to test these theoretical models, put constraints on the feedback mechanisms, identify the properties of the AGN host galaxies, and understand the accretion processes onto SMBHs and their fueling mechanism. X-ray surveys allow us to identify AGN activity without contamination from the emission of the host galaxy, i.e., therefore efficiently detecting even low luminosity AGNs. In the current era of deep and wide-area X-ray surveys with extensive spectroscopic follow-up, measurements of the three-dimensional (3D) clustering of AGN in various redshift ranges are emerging (\\citealt{coil_georgakakis_2009}; \\citealt{gilli_daddi_2005}; \\citealt{yang_mushotzky_2006}). However, our knowledge of AGN clustering in the low redshift universe $(z \\lesssim 0.4$) is still poor, except for optically selected type II AGNs (\\citealt{wake_miller_2004}; \\citealt{li_kauffmann_2006}). This is due to the lack of observable comoving volume and the low number density of low-$z$ AGNs. Exceptionally large survey areas with good sensitivity are needed to acquire a sufficiently large number of objects for clustering measurements. To date, the {\\em ROSAT} All-Sky Survey (RASS; \\citealt{voges_aschenbach_1999}) is the most sensitive survey to have mapped the entire sky in X-rays. Surveys with modern higher-sensitivity X-ray observatories such as {\\em XMM-Newton} and {\\em Chandra} cover much smaller areas of the sky (area: $\\sim$0.1-10 deg$^2$). Therefore, the available comoving volume from these deeper data sets is not sufficient to accurately measure AGN clustering at low redshifts $(z \\lesssim 0.4$). Serendipitous surveys, such as extended ChAMP (\\citealt{covey_agueros_2008}) and 2XMM (\\citealt{watson_schroeder_2009}), cover larger areas ($\\sim$33 deg$^2$ and $\\sim$360 deg$^2$, respectively). However, the large variations in sensitivity between different observations and the non-contiguous sky coverage make serendipitous surveys unsuitable for wide-area clustering measurements. A few studies have attempted to measure the auto-correlation function (ACF) of RASS-based AGN samples. Two-dimensional (2D) angular correlation functions (\\citealt{akylas_georgantopoulos_2000}) do not require redshift measurements for each AGN, but the projection heavily dilutes the clustering signal. Furthermore, the deprojection of the angular clustering to the 3D correlation function is subject to uncertainties in the redshift distribution, which can be substantial. Direct measurements of the 3D RASS-AGN ACF with spectroscopic redshift measurements have been made by \\cite{mullis_henry_2004} and \\cite{grazian_negrello_2004} with a few hundred AGNs, respectively, providing clustering measurements that have large statistical uncertainties caused by the relatively small sample size. \\cite{anderson_voges_2003, anderson_margon_2007} positionally cross-correlated RASS sources with spectroscopic data available from the Sloan Digital Sky Survey (SDSS). This dramatically increased the number of RASS-AGNs with spectroscopic redshift measurements, which we use here to provide significant improvements in the measurement of AGN clustering at low redshift. Furthermore, the availability of spectroscopic redshifts for large samples of SDSS galaxies in the same volume allows us to use an alternative approach to infer the clustering of AGN using calculations of the AGN--galaxy cross-correlation function (CCF). This approach uses much larger samples of AGN--galaxy pairs and hence significantly reduces the uncertainties in the spatial correlation function compared with direct measurements of the AGN ACF. Furthermore, the use of a CCF avoids the problem that we have to correct for the complex angular dependences of limiting sensitivity in the X-ray sample. We therefore have initiated a program to investigate the clustering properties of low redshift ($z\\sim 0.25$) RASS-AGNs through measurements of the CCF of these AGNs with SDSS Luminous Red Galaxies (LRGs). In this study, we chose LRGs as the corresponding galaxy sample because they have a significant overlap in redshift range with our X-ray sample (details are described later). In this first paper of a series, we explain the data selection, as well as the calculation of the CCF and the inferred RASS-AGN ACF. We also investigate the X-ray luminosity dependence of the clustering properties and biases. In a follow-up paper (T. Miyaji et al., in preparation), we will focus on applying the halo occupation distribution (HOD) model to the calculated CCF between RASS-AGNs and LRGs. The paper is organized as follows. In Section~2, we describe the construction and properties of the LRG and X-ray AGN samples in details. All essential steps to measure the CCF, compute the ACF via the CCF, and estimate errors are explained in Section~3. The results of the clustering measurements for the different X-ray AGN samples and their luminosity dependence are given in Section~4. We discuss these results in Section~5 in the context of other studies and conclude in Section~6. Throughout the paper, all distances, luminosities, and absolute magnitudes are measured in comoving coordinates and given in units of $h^{-1}$\\,Mpc, where $h= H_{\\rm 0}/100$\\,km\\,s$^{-1}$, unless otherwise stated. We use a cosmology of $\\Omega_{\\rm M} = 0.3$ and $\\Omega_{\\rm \\Lambda} = 0.7$ (\\citealt{spergel_verde_2003}). We use AB magnitudes throughout the paper. All uncertainties represent a 1$\\sigma$ (68.3\\%) confidence interval unless otherwise stated. ", "conclusions": "Using cross-correlation measurements between RASS-AGNs and SDSS LRGs, we measure the projected cross-correlation $w_p(r_p)$, and from this we derive the real-space auto-correlation function $\\xi(r)$ of X-ray selected AGNs with broad optical emission lines at $\\langle$$z$$\\rangle$$=0.25$ with an unprecedented precision. Using the RASS/SDSS cross-identification sample by \\cite{anderson_margon_2007} results in the largest sample of X-ray selected AGNs ($n=1552$) used for clustering studies, covering an area of 5468 deg$^2$. We detect a clustering signal at the $\\sim$11$\\sigma$ level and find a correlation length of $r_0=4.28^{+0.44}_{-0.54}$ $h^{-1}$ Mpc and $\\gamma =1.67^{+0.13}_{-0.12}$, fitting on scales of $r_p=0.3-15$ $h^{-1}$ Mpc. We investigate the luminosity dependence of clustering using low and high X-ray luminosity subsamples defined using the common AGN/QSO dividing line of $L_{\\rm X}^{0.5-10\\,{\\rm keV}} = 10^{44}$\\,erg\\,s$^{-1}$. We detect an X-ray luminosity dependence of the clustering signal at the $\\sim$2.5$\\sigma$ level, with the brighter sample being more clustered. This is contrasted with clustering measurements of low luminosity X-ray AGNs and optically selected QSOs at redshift $z>0.5$, which suggest that low X-ray luminosity AGNs are more strongly clustered and reside in more massive dark matter halos than their high X-ray luminosity counterparts. Possible explanations are that (1) different mechanisms trigger AGN activities at different redshifts and/or halo masses or (2) at all redshift, low luminosity AGNs and the brightest QSOs reside in red galaxies, while most intermediate luminosity AGNs/QSOs are hosted in blue galaxies. Our low $L_{\\rm X}$ RASS-AGN sample exhibits a similar clustering amplitude as blue, star-forming galaxies at similar redshifts, while the high $L_{\\rm X}$ RASS-AGN sample clusters like red galaxies. The total RASS-AGN sample is likely dominated by blue host galaxies but includes a fair fraction of red host galaxies. We show that the auto-correlation function derived from cross-correlation measurements of $\\sim$1500 AGNs and $\\sim$50000 LRGs in an area of $\\sim$6000 deg$^2$ constrains the auto-correlation function at the 10\\% level. Although the RASS is the most sensitive X-ray survey ever performed, only the most luminous and X-ray unabsorbed (soft) AGNs were detected. The upcoming all-sky {\\em eROSITA} mission (\\citealt{predehl_andritschke_2007}) with its {\\em XMM-Newton}-like soft and hard energy high sensitivity should detect $\\sim$200,000 AGNs with much lower X-ray luminosities. The data generated by the mission will allow one to compute clustering measurements with a significantly higher accuracy. To draw meaningful conclusions from AGN clustering measurements using AGN samples selected at different wavelengths and to address the evolution and luminosity dependence of AGN clustering, lower uncertainties on the detected clustering signal are required. We describe and successfully show in this paper that the cross-correlation function can be used to not only measure the relative clustering of AGN/galaxy samples but also to determine the auto-correlation function with relatively low uncertainties." }, "1002/1002.1776_arXiv.txt": { "abstract": "{ A sample of large northern Spitzer Infrared Nearby Galaxies Survey (SINGS) galaxies was observed with the Westerbork Synthesis Radio Telescope (WSRT) at 1300 -- 1760 MHz. In Paper II of this series, we described sensitive observations of the linearly polarized radio continuum emission in this WSRT-SINGS galaxy sample. Large-scale magnetic field structures of two basic types are found: (a) disk fields with a spiral topology in all detected targets; and (b) circumnuclear, bipolar outflow fields in a subset. Here we explore the systematic patterns of azimuthal modulation of both the Faraday depth and the polarized intensity and their variation with galaxy inclination. A self-consistent and fully general model for both the locations of net polarized emissivity at 1 -- 2 GHz frequencies and the global magnetic field topology of nearby galaxies emerges. Net polarized emissivity is concentrated into two zones located above and below the galaxy mid-plane, with the back-side zone suffering substantial depolarization (by a factor of 4 -- 5) relative to the front-side zone in its propagation through the turbulent mid-plane. The field topology which characterizes the thick-disk emission zone, is in all cases an axisymmetric spiral with a quadrupole dependance on height above the mid-plane. The front-side emission is affected by only mild dispersion (10's of rad~m$^{-2}$) from the thermal plasma in the galaxy halo, while the back-side emission is affected by additional strong dispersion (100's of rad~m$^{-2}$) from an axi-symmetric spiral field in the galaxy mid-plane. The field topology in the upper halo of galaxies is a mixture of two distinct types: a simple extension of the axisymmetric spiral quadrupole field of the thick disk and a radially directed dipole field. The dipole component might be a manifestation of (1) a circumnuclear, bipolar outflow, (2) an {\\it in situ} generated dipole field, or (3) evidence of a non-stationary global halo. } ", "introduction": "} The magnetic fields in spiral galaxies are an important component, but their basic three dimensional topology remains largely unknown. Two of their main characteristics are however, known. First, the fields in relatively face-on spiral galaxies are seen to follow the spiral pattern traced in the optical morphology. In the handful of more edge-on galaxies that have been imaged to date, the field distributions are seen to extend into the halo regions, and have a characteristic {\\sf X}-shaped morphology \\citep[eg.][]{heesen_etal_2009}. Apart from these basic properties, the details of the magnetic field topology are unknown. Observations of polarized flux, polarization vector orientations, and Faraday rotation measures all provide information about the magnetic field associated with different electron populations and at different projections with respect to the line of sight. Synchrotron emission originates in ultrarelativistic electrons spiralling around magnetic field lines, is beamed in the direction of motion of the electron, and is polarized perpendicular to the orientation of the field line. Polarized synchrotron radiation and polarization vector orientation are thus direct tracers of the magnetic fields perpendicular to the line-of-sight (LOS), $B_{\\perp}$, within the region where both ordered magnetic fields and relativistic electrons are maximized. The Faraday rotation measure (RM), or more generally the Faraday depth, $\\Phi$, that pertains to a given component of polarized emission, is sensitive to the integrated product of magnetic field component parallel to the LOS ($B_{\\parallel}$) and the thermal electron density in the foreground of a polarized emission component: \\begin{equation} \\Phi\\,\\propto\\,\\int_{\\mathrm{source}}^{\\mathrm{telescope}}n_e\\vec{B}{\\cdot}d\\vec{l}. \\end{equation} The Faraday depth is defined to be positive when $\\vec{B}$ points toward the observer, and negative when $\\vec{B}$ points away. When assessing the magnetic field geometry traced by these observational characteristics, it is essential to keep in mind that the observable attributes may originate in distinct regions of space. The classical Faraday rotation measure (RM) is an observable quantity derived from the polarization angle difference(s) $\\Delta\\chi$ between two (or more) frequency bands as $RM=\\Delta\\chi/(\\lambda_1^2-\\lambda_2^2)$. The empirically determined RM is only equivalent to the Faraday depth $\\Phi$ for a simple background emitter plus foreground dispersive screen geometry. Polarized emission can become depolarized in a number of ways: beam depolarization can arise because the spatial resolution element is large relative to the size of significant variations in the field orientation or the thermal electron content, while Faraday depolarization can arise because synchrotron emission and Faraday rotation take place in the same extended volume along the LOS. Polarized emission from different locations (either separated spatially or along the LOS) is affected by different amounts of Faraday rotation, such that at a given wavelength there may be orthogonal polarization angles that cancel, yielding no net polarization at that wavelength. Beam depolarization can be circumvented in principle by using higher angular resolution, although the brightness sensitivity may then be insufficient to detect the extended emission at all. Faraday depolarization can be circumvented in principle by achieving a sufficiently complete sampling of the $\\lambda^2$ measurement space (relevant for measurements of RM and $\\Phi$), since cancellation effects are confined to discrete wavelengths or ranges of wavelength. All of these observables can be used to constrain the likely magnetic field topology in the galaxies observed. In a previous paper (\\citet{heald_etal_2009}, hereafter Paper II), we presented our polarimetric results for a large sample of nearby galaxies observed to a comparable sensitivity limit of about 10~$\\mu$Jy beam$^{-1}$ RMS. In this paper, we begin by briefly summarizing the observations and data reduction steps of Paper II in Sect.~\\ref{section:reductions}. Trends noted in the data are described in Sect.~\\ref{section:trends}. In Sect.~\\ref{section:Bdist}, we then explore how particular magnetic field geometries might relate to the observations. We conclude the paper in Sect.~\\ref{section:disc}. ", "conclusions": "} The observational parameters and data reduction techniques of the WSRT-SINGS survey were presented in detail both by \\citet{braun_etal_2007}, and specifically regarding the polarization data in Paper II. Here we recap the most important details. For more information, the reader is referred to Sect.~2 and Appendix~A of Paper II. The data used in this analysis were obtained using the Westerbork Synthesis Radio Telescope (WSRT). Two observing bands were used: of $1300-1432$ MHz and $1631-1763$ MHz (centered on 22- and 18-cm, respectively), there being 512 channels in each band and in all four polarization products. Each galaxy in the WSRT-SINGS sample (refer to Paper II) was observed for 12~hr in the Maxi-short configuration of the WSRT. During each 12~hr synthesis, the observing frequency was switched between the two bands every 5~min. This provided an effective observation time of 6~hr per band, and good $uv$ coverage in both bands. The data for both bands were analyzed using the Rotation Measure Synthesis (RM-Synthesis) technique \\citep[][see also Paper II]{brentjens_debruyn_2005}. This provides the possible reconstruction of the intrinsic polarization vectors along each LOS, within the constraints set by the observing frequencies. The output of the RM-Synthesis procedure was deconvolved along the Faraday depth ($\\Phi$) axis, as described in Paper II. Polarized fluxes, polarization angles, and Faraday depths were extracted from these data and are discussed for each target galaxy in Paper II. In that paper, we also estimate the contribution to the RM from the Milky Way foreground using only background radio sources in the observed fields, rather than the target galaxies themselves. With this collection of data, we noted several patterns in the target galaxies, and that the basic patterns were common to the sample galaxies collectively. In this paper, we seek to explain these patterns using a common global magnetic field topology. } Our analysis of the projected three-dimensional magnetic field topologies presented in Sect.~\\ref{section:Bdist} and their predicted observable consequences for the azimuthal modulation, $B_{\\parallel}(\\phi)$ and $B_{\\perp}(\\phi)$, has provided a plausible explanation of the very general observed trends noted in Sect.~\\ref{section:trends}. A self-consistent scenario has emerged that accounts for the polarized intensity and its Faraday dispersion observed at GHz frequencies from galaxy disks. The detected polarized intensity is dominated by a zone of emissivity above the mid-plane on the side of the galaxy facing the observer (at a height of perhaps 5 -- 10 \\% of the disk radius). This thick disk emission arises in a region that is dominated by an axisymmetric spiral with an out-of-disk quadrupole topology, which is responsible for a distinctive modulation of $B_{\\perp}(\\phi)$ and its variation with galaxy inclination. This emission is affected by only a modest amount of Faraday dispersion, of a few tens of rad~m$^{-2}$, within the nearside halo of the galaxy in its subsequent propagation. For the majority of low to modest inclination galaxies ($\\le60^\\circ$), the dispersive foreground topology is consistent with an extension of the thick disk ASS quadrupole out to larger heights above the mid-plane (of perhaps 30\\% of the disk radius). The most highly inclined galaxies of our sample require an alternative halo field topology, in the form of a radially-dominated dipole, which yields a distinctive doubly-peaked modulation of $\\Phi(\\phi)$. It seems significant that in many or possibly all of the highly inclined galaxies of our sample there is evidence of a significant circum-nuclear outflow component to the polarized emission, in addition to that of the disk. This circum-nuclear component would quite naturally be expected to be associated with a dipole, rather than a quadrupole field, in view of the likely dominance of the $\\alpha^2$ over the $\\alpha\\omega$ dynamo mechanism at small galactic radii \\citep[e.g.][]{elstner_etal_1992}. This circum-nuclear dipole field would also be less likely to have any association with the spiral pitch angle of the disk given its origin. Because of the shallower roll-off with radius of a dipole compared to the quadrupole field (see Eq.~\\ref{eqn:dq}), a dipole component may come to dominate the halo field of the associated galaxy when both are present. \\citet{sokoloff_shukurov_1990} also argued that a $\\alpha\\omega$ dynamo operating in the halo would directly produce a dipole field. Non-stationary global halo models \\citep[e.g.][]{brandenburg_etal_1992} may also provide a natural explanation of the dipole signature on the largest scales. In addition to the bright polarized emission originating in the nearside, the corresponding rear-facing region of polarized emissivity of the thick disk can also be detected in relatively face-on galaxies if sufficient sensitivity is available. This emission is substantially weaker, by a factor of 4 -- 5, and consistent with depolarization caused by fluctuations in the magneto-ionic medium of the mid-plane on scales smaller than a pc. This fainter polarized component is affected by much greater Faraday dispersion, corresponding to both plus and minus 150 -- 200~rad~m$^{-2}$ in its propagation through the dense mid-plane plasma, as well as the near-side halo. These two maxima (one positive and one negative) in Faraday depth are aligned approximately with the major axes of each galaxy, and have approximately symmetric excursions about the Galactic foreground value. This pattern is consistent with the expectation for a simple planar ASS field in the galaxy disk. Future observations of nearby galaxy disks at frequencies below 200 MHz, such as with the upcoming LOFAR facility, will likely detect net polarized emission from even larger heights above the galaxy mid-plane and exclusively from regions unobstructed by the mid-plane in projection. A good indication for the predicted observables is given in Fig.~\\ref{figure:hsims} in which we present model integrations of the upper halo (out to 30\\% of the disk radius)." }, "1002/1002.4962_arXiv.txt": { "abstract": "In this paper, we analyze the full evolution, from a few days prior to the eruption to the initiation, and the final acceleration and propagation, of the CME that occurred on 2008 April 26 using the unprecedented high cadence and multi-wavelength observations by STEREO. There existed frequent filament activities and EUV jets prior to the CME eruption for a few days. These activities were probably caused by the magnetic reconnection in the lower atmosphere driven by photospheric convergence motions, which were evident in the sequence of magnetogram images from MDI (Michelson Doppler Imager) onboard SOHO. The slow low-layer magnetic reconnection may be responsible for the storage of magnetic free energy in the corona and the formation of a sigmoidal core field or a flux rope leading to the eventual eruption. The occurrence of EUV brightenings in the sigmoidal core field prior to the rise of the flux rope implies that the eruption was triggered by the inner tether-cutting reconnection, but not the external breakout reconnection. During the period of impulsive acceleration, the time profile of the CME acceleration in the inner corona is found to be consistent with the time profile of the reconnection electric field inferred from the footpoint separation and the RHESSI 15-25 keV HXR flux curve of the associated flare. The full evolution of this CME can be described in four distinct phases: the build-up phase, initiation phase, main acceleration phase, and propagation phase. The physical properties and the transition between these phases are discussed, in an attempt to provide a global picture of CME dynamic evolution. ", "introduction": "Coronal mass ejections (CMEs) are large-scale activities releasing a vast amount of plasma and solar energetic particles (SEPs) into the outer space \\citep{gos93,webb94}. These plasma and SEPs can propagate into the magnetosphere near the Earth and severely affect space-based modern technological systems, especially during the solar maximum \\citep{smart89}. Solar physicists have been pursuing what happens prior to the CME initiation and how CMEs are initiated. Various observational signatures, including magnetic cancellation, magnetic flux emergence, sigmoids, and filament activities are regarded as the significant precursors of CME eruptions \\citep{martin98,canfield00,wang06,gibson2006}. The common nature of these signatures are magnetic free energy build-up in the corona. As a consequence of the energy build-up, the coronal magnetic fields may explosively erupt once a trigger leads to the loss of equilibrium \\citep{forbes06}. However, there is no consensus so far on the exact trigger mechanism. The MHD instability model suggests that the eruption of CMEs that have a flux rope morphology is probably caused by the kink and/or torus instability when the winding of the field lines exceeds a critical value \\citep{sturrock01,linker01,fan04,rust05,totok05,gibson06}. In the tether-cutting model, the magnetic reconnection that occurs close to the polarity inversion line plays the role of weakening the constraining tension force of the overlying field, and results in the rise of the sigmoid-shaped core field and subsequently the runaway eruption \\citep{moore80,moore01,sturrock89,liu07,sterling07}. The same tension reduction mechanism holds for the flux emergence model suggested by Chen et al. (2000) in which the magnetic reconnection occurs between the emerging field and the background field. In the breakout model proposed by Antiochos et al. (1999), the overlying magnetic field constraining the sheared core field is removed through external magnetic reconnection, which leads to the CME eruption. Other authors have proposed that the injection of poloidal magnetic flux (of sub-photospheric origin) in the flux rope can cause a CME to take off \\citep{chenj00,chenj03,krall01}. More details about CME initiation mechanisms can be found in the reviews (e.g., Forbes 2000, 2006; Gopalswamy 2003; Chen 2008; Schrijver 2009). One key aspect of understanding CME eruption is of understanding the relationship between CMEs and flares, which itself has been a long-standing elusive issue for decades \\citep{kahler92,gos93,hund99}. Zhang et al. (2001, 2004) proposed three phases of CME kinematic evolution: the initiation phase, impulsive acceleration phase, and propagation phase, which are tightly associated with the three phases of the associated flare: the pre-flare phase, flare rise phase, and flare decay phase, respectively (see also, Burkepile et al. 2004; Vr\\v{s}nak et al. 2005b). The temporal correlation between CME acceleration and flare HXR flux was studied by Qiu et al. (2004) and Temmer et al. (2008). In the standard CME-flare model, the flare ribbons separate in the chromosphere during the CME impulsive acceleration phase because of continuous magnetic field reconnection. The reconnection rate can be calculated in terms of flare ribbon separation speed and the line-of-sight component of magnetic fields \\citep{forbes84,poletto86,forbes00,qiu02}. Qiu et al. (2004) compared the reconnection rate with the acceleration of the filament/CME and found a similarity between them. It was also found that the total reconnection flux is proportional to the maximum speed of CMEs (e.g., Qiu et al. 2005). In addition, Liu et al. (2009) found that the spectral index of X-ray emission of flares is strongly anti-correlated with the reconnection electric field. All of these suggest that CMEs and the associated flares, during the impulsive energy-release phase in particular, are driven by the same physical process in the lower corona, presumably via magnetic reconnection \\citep{lin00,priest02,vr04b,zhang06,mari07,temmer08}. As a matter of fact, most previous studies concerning CMEs address only certain specific phases of CME evolution, while very few are for the full evolution cycle from the build-up phase (tens of hours prior to the CME initiation), throughout the initiation phase and acceleration phase, and to the propagation phase. Therefore, it is useful to make a complete observation to investigate the full CME evolution. It is also of particular interest to study the variation of magnetic topology involved in different evolution phases, which shall shed light on possible initiation mechanisms of CMEs. The unique data of high cadence and full coverage acquired by SECCHI (Sun Earth Connection Coronal and Heliospheric Investigation; Howard et al. 2008) instruments onboard STEREO (Solar Terrestrial Relations Observatory; Kaiser et al. 2008) spacecraft provide us the opportunity to make such a study. In this paper, we investigate the full evolution of the CME on 2008 April 26 which was well observed by STEREO. In $\\S2$, we describe the instruments and the data. Our analysis and results are shown in $\\S3$ and $\\S4$. In $\\S5$, a schematic model is proposed to explain the full evolution of the CME, followed by discussions and conclusions in $\\S6$. ", "conclusions": "The STEREO observations provide an unprecedented opportunity to investigate solar eruptions. In this paper, we presented multi-wavelength observations of the flare-associated CME that occurred on 2008 April 26. We had studied its evolution for a long period and discussed the full evolution in a four-phase scenario: the build-up phase, initiation phase, main acceleration phase, and propagation phase. During the build-up phase, the active filaments, instantaneous EUV jets, and a reversed-S sigmoid structure were observed. All the features were physically related to the persistent slow magnetic reconnection in the solar lower atmosphere, which was manifested as photospheric magnetic cancellation. Before the eruption, there was a long period of reconnection occurring in the lower layers resulting in the transferring and accumulation of magnetic free energy, as well as the formation of a magnetic structure favorable for eruption, i.e., the sigmoid structure in this event. Different from the process of flux cancellation, the emerging of magnetic flux may also play an important role in transferring magnetic free energy from the sub-photosphere into the corona \\citep{tian08,Archontis09}. MacNeice et al. (2004) showed that, as the magnetic field shear increases, the magnetic free energy is continuously accumulated. They also proposed that such a quasi-static energy accumulation phase is necessary for any fast CME eruption. The magnetic field shear can be caused by the convergence motion of opposite magnetic fluxes \\citep{titov08}. Using a nonlinear force-free field extrapolation, it was recently found that the accumulated magnetic free energy increases with time prior to the eruption \\citep{thalmann08,guo08,jing09}. Therefore, the build-up phase accumulates the sufficient magnetic free energy for the eventual initiation and the final eruption. In general, the initiation phase of a CME eruption is characterized by a slow rise of the CME flux rope. In the present event, the EUV emission started to brighten in the core part of the reversed-S sigmoid configuration, which implied that slow magnetic reconnection was taking place there. As the field lines in the reversed-S sigmoid configuration continually reconnected, the CME flux rope rose slowly for about 20 minutes. This inner core magnetic reconnection prior to the eruption, combined with the facts of the bipolar magnetic structure in the active region and the absence of remote brightenings, seems to rule out the breakout model as the trigging mechanism of this event. Instead, we think that this eruption is well consistent with the tether-cutting initiation model. We also investigated the kinematics of this CME and found that its acceleration was well correlated with the HXR flux of the associated flare and the magnetic reconnection rate (see also, Zhang et al. 2004; Qiu et al. 2004; Temmer et al. 2008). It suggests that the main acceleration phase of the CME in the inner corona is likely caused by the fast runaway magnetic reconnection. Later on, the CME propagated with almost a constant velocity in the outer corona (see also, Zhang et al. 2001, Gallagher et al. 2003). However, for CMEs associated with a long decay flare, even though the fast magnetic reconnection ceases, a positive post-impulsive-phase acceleration may continue to exist after the impulsive acceleration phase \\citep{cheng09}. In general, these observational results are consistent with the standard CME-flare model. We think that the schematic model that comprises of the four phases proposed in this paper can be applied to most CME events. However, owing to different physical circumstances under which CMEs occur, individual events may have their own characteristics. In particular, there may be various manifestations for the build-up phase and the initiation phase, as mentioned above. We look forward to more observations in the coming years to study a variety of CME events, in order to fully understand the full evolution cycle of CMEs, including energy build-up, initiation, impulsive acceleration and subsequent propagation in the interplanetary space." }, "1002/1002.1892_arXiv.txt": { "abstract": "We present an ultradeep $K_S$-band image that covers $0.5\\times0.5$~deg$^2$ centered on the Great Observatories Origins Deep Survey-North (GOODS-N). The image reaches % a 5 $\\sigma$ depth of $K_{S,\\rm AB}=24.45$ in the GOODS-N region, which is as deep as the GOODS-N \\emph{Spitzer} Infrared Array Camera (IRAC) 3.6~$\\mu$m image. We present a new method of constructing IRAC catalogs that uses the higher spatial resolution $K_S$ image and catalog as priors and iteratively subtracts fluxes from the IRAC images to estimate the IRAC fluxes. Our iterative method is different from the $\\chi^2$ approach adopted by other groups. We verified our results using data taken in two different epochs of observations, as well as by comparing our colors with the colors of stars and with the colors derived from model spectral energy distributions (SEDs) of galaxies at various redshifts. We make available to the community our WIRCam $K_S$-band image and catalog (94951 objects in 0.25~deg$^2$), the Interactive Data Language (IDL) pipeline used for reducing the WIRCam images, and our IRAC 3.6 to 8.0~$\\mu$m catalog (16950 objects in 0.06~deg$^2$ at 3.6~$\\mu$m). With this improved $K_S$ and IRAC catalog and a large spectroscopic sample from our previous work, we study the color-magnitude and color-color diagrams of galaxies. We compare the effectiveness of using $K_S$ and IRAC colors to select active galactic nuclei (AGNs) and galaxies at various redshifts. We also study a color selection of $z=0.65$--1.2 galaxies using the $K_S$, 3.6~$\\mu$m, and 4.5~$\\mu$m bands. ", "introduction": "Deep near-infrared (NIR) imaging is essential for understanding galaxy evolution at high redshifts. One of its key applications is the determination of galactic mass. NIR luminosity is the best tracer of galactic stellar mass because it is less affected by extinction and young stars. This role is traditionally played by deep ground-based $K$-band imaging \\citep[e.g.,][]{cowie94,kauffmann98,brinchmann00}. An important reason for this is that the $K$-correction at $K$-band is small across a broad range of redshifts and is relatively invariant against galaxy types compared to the other two NIR bands ($J$ and $H$, see, e.g., \\citealp{keenan09}). Recently, the high sensitivity of the \\emph{Spitzer} Infrared Array Camera (IRAC, \\citealp{fazio04}) adds additional powerful wavebands for this purpose \\citep[e.g.,][]{fontana06,perez08,elsner08}. However, ground-based $K$-band imaging remains important because of the ease of observations. In order to study galaxy formation and evolution, we have been carrying out deep NIR imaging in the Great Observatories Origins Deep Survey-North \\citep[GOODS-N,][]{giavalisco04}. In \\citet[hereafter B08]{barger08} we presented a deep $K_S$-band catalog in the GOODS-N obtained with the Wide-field InfraRed Camera (WIRCam) on the 3.6~m Canada-France-Hawaii Telescope (CFHT). The WIRCam image was used by \\citet{wang07} to study an optically faint submillimeter galaxy. The catalog was used by \\citet{cowie08} to construct a mass-selected sample of galaxies at $z<1.5$ with an AB magnitude limit of 23.4. We have since added more data to the WIRCam $K_S$ imaging and slightly improved the reduction. In this paper we present the new image and a new $K_S$ catalog. In order to achieve a more complete sampling of galaxy mass at high redshifts, we have also constructed a new 3.6 to 8.0~$\\mu$m catalog based on \\emph{Spitzer} IRAC images in the GOODS-N (M.\\ Dickinson et al., in preparation), which we also present here. Colors in the \\emph{Spitzer} IRAC bands are used to study the redshifts of galaxies, especially optically faint ones \\citep[e.g.,][]{pope06,wang07,devlin09,marsden09}. This is largely based on the rest-frame 1.6~$\\mu$m bump in galactic spectral energy distributions (SEDs), which is caused by the H$^-$ opacity minimum in the stellar photosphere \\citep{simpson99,sawicki02}. IRAC colors are also used for separating active galactic nuclei (AGNs) from galaxies (\\citealp{lacy04,stern05}, see also \\citealp{hatzim05,sajina05}). In B08 we studied the effectiveness of using IRAC colors to select AGNs and galaxies at various redshifts based on a large spectroscopic sample, and we pointed out the limitations in these selections. With our improved color measurements in the IRAC and $K_S$ bands, we revisit and discuss these issues. We describe the observations in \\S~\\ref{sec_observation}, the data reduction in \\S~\\ref{sec_reduction}, the reduction quality in \\S~\\ref{sec_quality}, and a $K_S$ selected catalog in \\S~\\ref{k_catalog}. We present in \\S~\\ref{sec_color} a new method of constructing IRAC catalogs based on the $K_S$-band image and catalog, and we evaluate the performance of our new method. In \\S~\\ref{sec_properties} we describe the general properties of the $K_S$ and IRAC catalog, including its overlap with other multiwavelength catalogs, the $K_S$ and IRAC color-magnitude diagrams and color-color diagrams. We summarize our results in \\S~\\ref{sec_summary}. The images and catalog data from this work will be available online. All fluxes in this paper are $f_\\nu$. All magnitudes are in AB system, unless otherwise stated, where an AB magnitude is defined as $\\rm AB=23.9-2.5\\log(\\mathrm{Jy})$. ", "conclusions": "\\label{sec_summary} We carried out ultradeep $K_S$ band imaging in the HDFN and its flanking field with WIRCam on the CFHT. The mosaic image has a size of $35\\arcmin \\times 35\\arcmin$, fully covering the GOODS-N field. With nearly 50~hr of integration, we reached a 90\\% completeness limit of 1.38~$\\mu$Jy ($\\rm AB=23.55$) in the 0.25~deg$^2$ region, and 0.94~$\\mu$Jy ($\\rm AB=23.96$) in the 0.06~deg$^2$ GOODS-N IRAC region. The faintest detected objects have $K_S$ fluxes of $\\sim0.2$~$\\mu$Jy ($\\rm AB = 25.65$). The photometry is uniform across the entire field, and the variation is well within 2\\%. In the GOODS-N region the astrometry is accurate to $0\\farcs03$ for compact objects with high S/N. The $K_S$ images and the associated catalogs are available through the internet. We used the $K_S$ image and catalog as priors to measure MIR magnitudes in the \\emph{Spitzer} IRAC 3.6 to 8.0~$\\mu$m bands with our REALCLEAN method. We verified the REALCLEAN results using data taken in two different epochs of observations, as well as by comparing our colors with the colors of stars and with the colors derived from model SEDs of galaxies at various redshifts. The REALCLEAN method appears to provide reasonably good photometry even on galaxies that have close neighbors within $2\\arcsec$. This also provides a relatively clean way to identify faint IRAC objects that are not detected in the $K_S$ image. We make the REALCLEAN IRAC catalogs for both $K_S$ detected and undetected objects available to the community. Additional to the REALCLEAN method, we also measured galactic colors by convolving the $K_S$ image with the IRAC PSFs and convolving the IRAC images with the $K_S$ PSF. The derived colors are consistent with the REALCLEAN results. We found our $K_S$ sample to be extremely powerful in studying the multiwavelength properties of high-redshift galaxies. The $K_S$ image detected most (if not all) of the X-ray, 24~$\\mu$m, and radio sources published in this field. On the other hand, only 40\\% of the \\emph{HST} ACS sources in the GOODS-N are detected at $K_S$ at the current depth. We studied various $K_S$ and IRAC color-magnitude and color-color diagrams. Because of the rest-frame 1.6~$\\mu$m bump in the SEDs of galaxies, galaxies show significant color evolution in the $K_S$ and IRAC bands from $z=0$ to $z<2$. We found that using IRAC colors to select galaxies at certain redshifts is extremely sensitive to the color measurements and that balancing selection completeness and contamination will require careful tuning of the selection criteria with a good spectroscopic sample. We discussed an effective selection of $z=0.65$--1.2 galaxies using the $K_S$ band and the two shorter IRAC bands that will remain available in the \\emph{Spitzer} warm phase mission. We also confirmed our previous studies (B08) that the IRAC AGN selections in \\citet{lacy04} and \\citet{stern05} are either incomplete at the low-luminosity end or suffer severely from galaxy contamination." }, "1002/1002.4479_arXiv.txt": { "abstract": "Our aim is to investigate whether the presence of baryons can have any significant influence on the properties of the local Hubble flow which has proved to be ``cold''. We use two cosmological zoom simulations in the standard $\\Lambda$CDM cosmology with the same set of initial conditions to study the formation of a local group-like system within a sphere of $\\sim 7\\,h^{-1}$ Mpc. The first one is a pure dark matter simulation (\\rundm) while a complete treatment of the physics of baryons is introduced in the second one (\\runb). A simple algorithm based on particles identity allows us to match haloes from the two runs. We found that galaxies identified in \\runb \\, and their corresponding dark matter haloes in \\rundm \\, have very similar spatial distributions and dynamical properties on large scales. Then, when analyzing the velocity field and the deviation from a pure Hubble flow in both simulations, namely when computing the dispersion of peculiar velocities of galaxies $\\sigma_*(R)$ and those of their corresponding dark matter haloes $\\sigma_{DM}(R)$ in \\rundm, we found no particular differences for distances $R=$1 to 8 Mpc from the local group mass center. This suggests that the presence of baryons have no noticeable impact on the global dynamical properties of the local Hubble flow within such distances. Then, the results indicate that the ``true'' $\\sigma_{*}(R)$ values can be estimated from the pure dark matter simulation with a mean error of 3 km/s when dark matter haloes are selected with maximum circular velocities of $V_c\\geq$30 km/s, corresponding to a population of dark matter haloes in \\runb \\, that host galaxies. By investigating the properties of the Hubble flow at distances $R\\sim$0.7 to 3 Mpc, we also found that the estimation of the total mass enclosed at the radius of the zero-velocity surface $R_0$, using the spherical infall model adapted to $\\Lambda$CDM, can be underestimated by at least 50\\%. ", "introduction": "The study of dynamical and photometric properties of galaxies in the local universe represents an ideal framework to confront predictions of the standard $\\Lambda$ cold dark matter ($\\Lambda$CDM) model with observations. In particular, an interesting feature of the dynamical properties of the local universe is that the Hubble flow is rather ``cold'', e.g. the dispersion in the peculiar velocities within the Hubble flow is quite small, namely $\\leq$ 100 km/s (Sandage \\& Tammann 1975; Giraud 1986; Schlegel et al. 1994; Ekholm et al. 2001, Karachentsev et al. 2003; Macci{\\`o}, Governato \\& Horellou 2005; Tikhonov \\& Klypin 2009). The presence of the dark energy was proposed as a possible explanation for the smoothness of the local Hubble flow, first argue by Baryshev et al. (2001), Chernin et al. (2001, 2004, 2007) and Teerikorpi et al. (2005) and supported by Macci{\\`o}, Governato \\& Horellou (2005) using a set of N-body simulations. However, Hoffman et al. (2007) compared results from their own simulations of CDM and $\\Lambda$CDM cosmologies with identical parameters, apart from the presence or not of the cosmological constant term. They claimed that no significant differences were noticed in the velocity flow around galaxies having properties similar to those observed in the neighborhood of the Milky Way. A similar conclusion was obtained more recently by Martinez-Vaquero et al. (2009) using a range of N-body simulations in different cold dark matter scenarios. They conclude that the main dynamical parameter that can affect the coldness of the flow is the relative isolation of the Local Group. Another approach was proposed by Peirani \\& de Freitas Pacheco (2006, 2008) who derived a velocity-distance relation by modifying the Lema\\^itre-Tolman model (TL) with the inclusion of a cosmological constant term. They found that this new relation, which describes the behavior of the Hubble flow near the central dominant objects ($\\leq$2-3 Mpc), in addition to that derived from the ``canonical'' model ($\\Omega_m=1$), provide equally acceptable fits to the existing available data. As a consequence, no robust conclusion about the effects of the cosmological constant on the dynamics of groups could be established. Moreover, Axenides \\& Perivolaropoulos (2002) studied the dark energy effects in the growth of matter fluctuations in a flat universe. They concluded that the dark energy can indeed cool the local Hubble flow but the parameters required for the predicted velocity dispersion to match the observed values are out by observations that constrain either the present dark energy density or the equation of state parameter w(=$P_x/\\varepsilon_x$). Other cosmological models were proposed to study the local Hubble flow. Dark matter simulations by Governato et al. (1997) for cosmological models with $\\Omega_m = 1$ or $\\Omega_m = 0.3$ are, according to these authors, unable to produce systems embedded in regions of ``cold'' flows, i.e., with 1-D dispersion velocities of approximately 40-50 km/s. Recently, Tikhonov et al. (2009a, 2009b) have compared the observed spectrum of minivoids in the local volume with the spectrum of minivoids determined from the simulations of CDM or WDM models. They found that model predictions and observations match very well provided that galaxies can only be hosted by dark matter haloes with circular velocities greater than 20 km/s (for WDM) and 35 km/s (for \\lcdm). They have also derived rms deviations from the Hubble flow which seem to be consistent with observational values. All those past numerical works have used collisionless N-body simulations to investigate the puzzle of coldness of the local Hubble flow and thus, the velocity dispersions have been derived by considering the dynamical properties of dark matter haloes supposed to host galaxies. However, it is nonetheless not obvious that positions and velocities of galaxies and their host dark matter haloes are identical. A first element of answer was given by Weinberg et al. (2008) who have studied the subhalo and galaxy populations in a galaxy group simulated either by a collisionless cosmological simulation or a hydrodynamics simulation with the same initial conditions. They found that positions and masses of large subhaloes are very similar in both runs while they can be different for low mass subhaloes whose orbits can be more easily modified by the host halo potential. Past works also suggest the existence of bias between the star and dark matter components (see for instance, Carlberg, Couchman \\& Thomas 1990; Zhao, Jing \\& B\\\"orner 2002; Sousbie et al. 2008). In the present work, we study possible effects of the presence of baryons in the dynamical properties of the local Hubble flow by using high resolution simulations that include most of the relevant physical processes that lead to the formation of galaxies inside dark matter haloes. This paper is organized as follows: in Sect. 2, we summarize the numerical methodology; in Sect. 3 we present results on the dispersions of peculiar velocities around the mean Hubble local flow derived from our numerical models and, finally, in Sect. 4, our main conclusions are given. ", "conclusions": "In the present paper, we have used cosmological N-body simulations with and without a complete treatment of the physics of baryons to study the formation of a local group type halo. We have first identified a LG system in a simulation with a low mass particle resolution, using standard criteria. Then, the best candidate has been re-simulated with higher resolution using the technique of zoom. This approach differs from the technique of constrained simulations used for instance in the CLUES project (Yepes et al. 2009), which include the correct motion and position of objects in a large volume such as the Local Supercluster, the Virgo and the Coma cluster and the Great Attractor (see also Lavaux 2010). We focus on the dynamical properties of the local Hubble flow and in particular, the influence of baryons. From the two runs, we built mock catalogs of either galaxies or dark matter haloes and computed their velocities as respect to the mass center of the central main pair (Milky Way-M31). We found that the dispersions of peculiar velocities around the mean local Hubble flow of galaxies in \\runb \\, and those of their corresponding dark matter haloes in \\rundm \\, are very close for distances $D=1$ to 8 Mpc. Moreover, similar $\\sigma_H$ values (with mean errors $\\leq 5$ km/s) are also obtained for samples of dark matter haloes extracted from \\rundm \\, with constraint of their maximum circular velocity. These results suggest that the global dynamical properties of the Hubble flow is not affected by the presence of baryons. Such a result is not very surprising since galaxies are expected to form in the gravitational potential well of dark matter haloes. Therefore, their properties are expected to be closely related to those of their host haloes. This statement is confirmed on large scales as shown in Figs. \\ref{fig_dist} and \\ref{fig_prob}. However, some discrepancies between positions and dynamical properties of galaxies and those of their corresponding haloes may raised in the case of low mass objects in denser environment. For instance, the presence of baryons is supposed to favor the survival of substructures inside massive haloes against tidal striping (see Duffy et al. 2010 and references therein) and can affect their distribution (Weinberg et al. 2008; Libeskind et al. 2010, Knebe et al. 2010). But these effects seem to have no particular consequences in the estimation of the global $\\sigma_H$ values. It should also be pointed out that \\runb \\ does not include any prescriptions of strong feedback such as galactic winds or Sedov shock waves that are expected to lead to a significant ejection of gas from the central region of forming galaxies. As a result, the fraction of baryons (in the form of stars and cold gas) in the central region of simulated massive galaxies such as the Milky Way or M31 is close to the cosmological baryonic fraction (see Table \\ref{table1}) whereas observations suggest lowest values (Fukugita, Hogan \\& Peebles 1998; Bell et al. 2003). However, no one has solved this \u201cmissing baryon problem\u201d yet and recent studies even suggest that theoretical sources of strong winds such as AGN and supernovae are insufficient to explain the missing baryons in Milky Way type galaxies (see for instance Anderson \\& Bregman 2010; Silk \\& Nusser 2010). Although the present study mainly focus on lower mass galaxies ($\\leq 10^{10} M_\\odot$) in which supernovae feedback is supposed to have a significant impact in their formation, the other sources of strong feedback should be taken into account as well. Then, a natural extension of this work would be to add these processes while increasing the mass resolution in the simulation (see for instance Ceverino \\& Klypin 2009) to see whether the conclusions presented here (especially the distribution of galaxies as respect to their corresponding dark matter haloes) are significantly affected of not. In good agreement with previous theoretical and numerical works, we found that the halo population in the pure dark matter run can be divided into two groups according to their maximum circular velocity $V_c$. There is indeed a critical value, $V_c\\sim 30-35$ km/s (corresponding to haloes with mass lower than $\\sim 7\\times 10^{9}\\,h^{-1}M_\\odot$ in \\rundm), below which the corresponding SPH haloes don't host any galaxy. Such a phenomenon may be explained by the fact that the cooling of gas is expected to be suppressed by feedback processes such as a UV-photoionisation in low mass haloes (see references in paragraph 2.3). As regards this point, we cannot exclude the fact that the numerical resolution used in this work may affect the results especially for low mass systems. However, haloes with mass of $\\sim 7\\times 10^{9}\\,h^{-1}M_\\odot$ are composed by $\\sim 800$ particles and this seems reasonable for the estimation of the thermodynamic properties of the gas component. Then, table \\ref{table3} indicates that the number of dark matter haloes in \\rundm \\ selected with $V_c\\geq 30$ km/s is actually quite close to the real number of galaxies selected \\runb. Therefore, one can use this single criteria to select dark matter haloes in a collisionless simulation to have an good estimation of $\\sigma_H$. The existence of a dark halo population in the Local Group can also have important consequences in understanding the discrepancies between the large number of subhaloes present in simulations but not observed (Kauffmann, White \\& Guiderdoni 1993; Moore et al. 1999; Klypin et al. 1999). This problem is beyond the scope of this paper and will be investigated in detail in a forthcoming paper. The velocity dispersion obtained from our fiducial model (e.g. from galaxies) is in quite good agreement with observational expectations. For instance, with added corrections for distance errors, we obtained $\\sigma_{*,cor}=90.4$ for $D=1-8$ Mpc, a value which tend to be even lower that the one derived from observational analysis, namely 99.3 km/s (Tikhonov \\& Klypin 2009) and with a reasonable number of galaxies considered. For distances from 0.7 to 3 Mpc, we also found a nice agreement between our simulated data and observational ones from Karachentsev et al. (2009). Then, these results suggest that there is no particular problem with the \\lcdm \\ model in reproducing the velocity dispersion derived from observational analysis. A similar conclusion was obtained from Martinez-Vaquero et al. (2009) who found that a non-negligible fraction of LG like objects simulated form various cosmological models present a $\\sigma_H$ value close to (and even smaller) than the observed value. Moreover, similar results have been also obtained from the analysis of Tikhonov \\& Klypin (2009). Since the present study has been limited to only one realization, due to the high computational cost when the physics of baryons is included in the simulations, general conclusions cannot be drawn. However, one of our main criteria to select LG candidates was their relative isolation (no massive galaxies within 3 Mpc). We stress that we found only one candidate (among 16 MW-M31 pair candidates) that was satisfying all the criteria. This raised the question whether the local group is located in a particular place of the universe or not. To finish, we have also compared our simulated data to observational ones for $R=0.7-3$ Mpc. The nice agreement between them allows us to test models, in particular the estimation of the total mass enclosed at the zero-velocity surface radius $R_0$ using the Lema\\^itre-Tolman model adapted to $\\Lambda$CDM. It appears that the estimation of the LG mass using this approach is underestimated by at least 50\\%. This is probably due to the fact that, on the one hand, the hypothesis of a spherical infall collapse is not valid anymore in this specific case. On the other hand, the mass enclosed within $R_0$ is not constant in the time, as assumed by the TL model (e.g. no mass accretion), and this effect may have an important influence in the estimation of the final total mass. \\vspace{1.0cm} \\noindent {\\bf Acknowledgment} \\noindent I acknowledge support from the ``Agence National de la recherche'' ANR-08-BLAN-0222-02. It is a pleasure to thank the anonymous referee for his/her useful comments which have significantly improved this paper. I warmly thanks C.\\,Alard, S.\\,Colombi, J.\\,Devriendt, R.\\,Gavazzi, M.\\,Hudson, Y.\\,Kakazu, I.\\,D.\\,Karachentsev, A.\\,Klypin, G.\\,Mamon, R.\\,Mohayaee, J.\\,A. de Freitas Pacheco, C.\\,Pichon, S.\\,Prunet, J.\\,Silk and T.\\,Sousbie for interesting discussions. I also thanks D.\\,Munro for freely distributing his Yorick programming language (available at \\texttt{http://yorick.sourceforge.net/}) which was used during the course of this work. This work was carried within the framework of the Horizon project (http://www.projet-horizon.fr)." }, "1002/1002.0636_arXiv.txt": { "abstract": "We identified a large sample of radio quasars, including those with complex radio morphology, from the Sloan Digital Sky Survey (SDSS) and the Faint Images of Radio Sky at Twenty-cm (FIRST). Using this sample, we inspect previous radio quasar samples for selection effects resulting from complex radio morphologies and adopting positional coincidence between radio and optical sources alone. We find that $13.0$\\% and $8.1$\\% radio quasars do not show a radio core within $1 \\farcs 2$ and $2^{''}$ of their optical position, and thus are missed in such samples. Radio flux is under-estimated by a factor of more than 2 for an additional $8.7\\%$ radio quasars. These missing radio extended quasars are more radio loud with a typical radio-to-optical flux ratio namely radio loudness $RL\\gtrsim 100$, and radio power $P\\gtrsim 10^{25}$ W Hz$^{-1}$. They account for more than one third of all quasars with $RL>100$. The color of radio extended quasars tends to be bluer than the radio compact quasars. This suggests that radio extended quasars are more radio powerful sources, e.g., Fanaroff-Riley type 2 (FR-II) sources, rather than the compact ones viewed at larger inclination angles. By comparison with the radio data from the NRAO VLA Sky Survey (NVSS), we find that for sources with total radio flux less than 3 mJy, low surface brightness components tend to be underestimated by FIRST, indicating that lobes in these faint radio sources are still missed. ", "introduction": "In past decades, we have witnessed a rapid growth in the number of radio selected Active Galactic Nuclei (AGNs) resulting from large and deep radio surveys such as the Faint Images of Radio Sky at Twenty centimeters (FIRST, Becker, White \\& Helfand 1995) and the NRAO VLA Sky Survey (NVSS, Condon et al. 1998) coupled with optical spectroscopy follow-ups or dedicated spectroscopic surveys such as the Two degree Field (2dF, Maddox et al. 1998; Boyle et al. 2001) and the Sloan Digital Sky Survey (SDSS, York et al. 2000; Stoughton et al. 2002). However, some fundamental issues regarding the origin of radio emission remain hotly debated: is the radio loudness dichotomy true or not (Kellermann et al. 1989; Miller, Peacock \\& Mead 1990; Hewett et al. 2001; Ivezic et al. 2002)? Are radio jets in radio quiet/intermediate quasars relativistic (Readhead et al. 1988; Wilson \\& Colbert 1995; Meier et al. 2001)? Which physical parameters control the large range of radio strength despite their great similarity in the SED at other wavelengths for various quasars (Barthel 1989; Urry \\& Padovani 1995; Jackson \\& Wall 1999; Boroson 2002; Aars et al. 2005)? Concerning the first question, the ambiguity is caused largely by various selection biases introduced by the survey limits, the incompleteness in the radio sample due to their complex morphologies or in the optical sample due to various color selections (Ivezic et al. 2002; Cirasuolo et al. 2003a; Best et al. 2005). The traditional radio surveys were carried out at shallow flux density limits (0.1-1 Jy) and primarily radio-loud quasars were detected (e.g., Bennett 1962; Smith \\& Spinrad 1976; Colla et al. 1972; Fanti et al. 1974; Large et al. 1981). Only the two latest radio surveys NVSS and FIRST have enough sensitivity and positional accuracy to allow for the detection of large number of radio intermediate and radio quiet quasars in conjunction with moderately deep, large area optical surveys such as 2dF and SDSS. However, by taking advantage of the accuracy in the position of radio sources, most authors constructed the radio quasar or quasar candidate sample using solely the position match between radio and optical sources (Gregg et al. 1996; White 1999; Lacy \\& Ridgway 2001; McMahon et al 2002; Richards et.al. 2002; Cirasuolo et al.2003b). This process introduces a bias against lobe-dominated radio quasars: either they were missed or their radio flux underestimated. White, Becker \\& Gregg (1999, 2000) argued that this incompleteness is not severe in FIRST Bright Quasar Survey (FBQS), which was selected by matching of optical counterparts within a $1\\farcs2$ position offset to the radio sources in the FIRST catalog. Using a matching radius of $2^{''}$, Ivezic et al. (2002) estimated that less than $10\\%$ SDSS-FIRST associations have complex radio morphology, and core-lobe and double-lobe sources together represent about only $5\\%$ in the radio quasars and galaxy sample. Using a novel technique, de Vries et al. (2006) constructed a Fanaroff-Riley type 2 (FR-II) quasar sample, and found that $27\\%$ FR-II quasars do not show cores at the FIRST flux limit. These authors also compared the emission line properties and optical colors of these FR II quasars with radio quiet quasars. It should be noted that these missing quasars are not random, but are all extended sources and tend to be more radio loud (Falcke et al 1996; Ivezic et al. 2002; Best et al 2005). As a result, the statistical properties of the sample, such as radio loudness and radio luminosity distribution, will be affected by this selection effect. In this paper we study in detail of selection effects in the SDSS radio quasar samples. We identified from SDSS and FIRST a large sample of radio quasars, including those with complex radio morphology. Besides using positional coincidence as a primary selection criterion, we manually examined the FIRST images for all of the candidates with extended radio morphology. Through this less efficient process, we obtained a sample of $3641$ spectroscopically confirmed quasars with secure radio identification. A detail comparison of this sample with other radio quasars sample is given. Various selection effects are quantified. Throughout this paper, we will adopt a concordant cosmology with $H_0 = 70$ km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m = 0.3$, and $\\Omega_{\\Lambda}= 0.7$. ", "conclusions": "We have constructed a relatively unbiased large radio quasar sample using the SDSS quasar catalog (Schneider et al. 2005), and the FIRST catalog and images. Apart from positional coincidence of radio sources within $2^{''}$ of quasars, we also identify the radio counterpart of quasars with complex radio morphologies such as lobe -dominated quasars by visual inspection of their radio images. We find that using the positional coincidence alone will miss $\\sim 8\\%$ radio counterpart that do not show radio core at FIRST flux limit of 1 mJy, and under-estimates the radio flux by a factor of more than two in another $\\sim 9\\%$ objects. By comparing the radio flux from FIRST survey with that from NVSS, we found that lobes in weak radio sources tend to be missed in this sample. So these numbers are only lower limits. Quasars with extended radio emission show both larger radio powers and radio loudness, and appear somewhat bluer than radio compact quasars despite their indistinguishable optical luminosity. As such, the radio extended quasars account for nearly one third of radio loud quasars at log$(RL)>2.2$. Naturally, including the extended emission and weak core radio sources increases the fraction of the radio loud objects and the significance of radio loud peak in the distribution of Hough \\& Readhead radio loudness. Our results in the first glance are not consistent with the simple unification scheme in which radio compact quasars are extended ones viewed along radio jet, for which the relativistic beaming enhances the core radio emission and as such the total radio power (Wills \\& Browne 1986; Hough \\& Readhead 1988; Barthel 1989; Falcke et al. 2004) when the projection effect would make the apparent size smaller. Within such a scheme, the unresolved core is enhanced because of the beaming effect. That the lobe-dominated quasars are more luminous in radio seems to contradict this model. However, there are at least two selection effects that make the average radio power in the core-dominated sources smaller. First, as we showed in the last section that the peak brightness of radio lobe component is correlated with radio power, and the FIRST will not be able to detect the lobe component in less radio power sources, especially at high redshift (see also Fig.\\ref {lobe-distrib}). Second, if most of core-dominated quasars are intrinsically radio weaker (Wang et al. 2006) and if the radio luminosity function of them is steep, their average apparent luminosity can be lower even if the radio power had been boosted. The relative number density of the extended and compact sources also suggests that majority of compact sources are either of beamed, intrinsically much weak radio quasars or intrinsically compact radio sources, such as CSS (Compact Steep Spectrum Objects) or GPS (GigaHertz Peaked Sources). With typical Lorentz factor of 10-15 for jets in FR-II radio quasars, the boosted emission can be viewed in only relatively small fraction of solid angle between the line of sight and the jet, $\\theta <7^{\\circ} $, whereas at other angles the core is weakened due to Doppler effect. If the power of the un-beamed radio cores is near 0.005 of that of the lobes, as determined for 3CR sources (Urry \\& Padovani 1995), and if the intrinsic radio loudness distribution following Cirasuolo et al. (2003b) , and with the luminosity distribution as DR3 quasar sample, we can estimate that $\\sim 62\\%$ beamed intrinsic radio quiet sources could be detected by FIRST. Since GPS and CSS are all powerful radio sources, it is likely most of these compact quasars are beamed radio intermediate quasars as proposed by Falcke et al. (1996) Our results suggest that lobe-dominated sources are not particularly reddened, in agreement with the finding by de Vries et al. (2005). This is valid even for quasars without detectable radio core at flux down to the FIRST limit. The line of sight does not intercept the dusty torus in those lobe-dominated quasars. However, Backer et al. (1997) found that most lobe-dominated quasars in the Molonglo quasars are reddened by $A_V\\simeq 2-4$, and CSSs are most reddened. It should be pointed out that quasars reddened by this amount cannot be found in the `color' selected sample, in particularly at high redshift due to strong attenuation in ultraviolet. The slightly bluer color for lobe-dominated quasars in this sample could be due to inclusion of CSS-like objects in the core-dominant objects or due to a selection effect by which reddened \"extended\" radio quasars are lost. If Baker (1997) is correct, we might miss a large number of heavily reddened lobe-dominated quasars. Although most such quasars are likely below the magnitude limit of spectroscopic quasar sample, in principle the FIRST selected sample is able to detect some of such reddened quasars, particularly in the low redshift if a weak core is present. We look at the spectroscopic sources that selected as FIRST sources only, and find that the FIRST-only selected spectroscopic sources are indeed much redder. But the fraction of quasars with extended lobes in FIRST-only sources is very low($\\sim 0.2\\%$) probably due to large extinction. Thus it is not conclusive whether a large number of such reddened lobe-dominated quasars do exist." }, "1002/1002.2069_arXiv.txt": { "abstract": "We report on the discovery of a new heavily polluted white dwarf. The DAZ white dwarf GALEX~J193156.8+011745 was identified in a joint {\\it GALEX}/GSC survey of ultraviolet-excess objects. Optical spectra obtained at ESO NTT show strong absorption lines of magnesium and silicon and a detailed abundance analysis based on VLT-Kueyen UVES spectra reveal super-solar abundances of silicon and magnesium, and near-solar abundances of oxygen, calcium, and iron. The overall abundance pattern bears the signature of ongoing accretion onto the white dwarf atmosphere. The infrared spectral energy distribution shows an excess in the H and K bands likely associated with the accretion source. ", "introduction": "The star GALEX~J193156.8+011745 (hereafter GALEX~J1931+0117) is a hydrogen-rich white dwarf recently identified in a joint ultraviolet/optical survey \\citep{ven2010} based on the {\\it Galaxy Evolution Explorer} ({\\it GALEX}) all-sky survey\\footnote{See \\citet{mor2007} for a description of the data products.} and the GSC2.3.2 catalogue. The Third U.S. Naval Observatory CCD Astrograph Catalog\\footnote{Accessed at VizieR \\citep{och2000}.} locates the star at R.A.(2000)$=19\\ 31\\ 56.933$, Dec.(2000)$=+01\\ 17\\ 44.13$ with a proper motion of $\\mu_\\alpha\\cos{\\delta}=-77.6\\pm6.4$ and $\\mu_\\delta=-6.1\\pm6.1$ mas~yr$^{-1}$. The distance modulus implies that the star is nearby ($\\approx55$ pc) and that it is imbedded in the Galactic plane ($|z|\\approx 8$ pc). Based on echelle spectroscopy we show that GALEX~J1931+0117 is a DAZ white dwarf and is one of the most extreme cases of externally polluted white dwarf atmospheres. The DAZ white dwarf atmospheres are hydrogen-dominated with heavy element abundances ranging from nearly solar down to detection limits several orders of magnitude below solar \\citep{zuc2003,koe2005}. The DAZ and, more particularly, the DZ phenomena are interpreted as evidence of on-going accretion of circumstellar material onto white dwarf atmospheres \\citep[e.g.,][]{kil2007}. The source of the material has been ascribed to comets \\citep{alc1986} and more recently to tidally-disrupted asteroids \\citep[see][]{jur2008}, although close, low-mass companions in post-common envelope binaries are also a known source of material \\citep{deb2006,kaw2008}. The observational evidence in favour of the accretion of debris material resides in (1) the inferred composition of the accretion flow \\citep{zuc2007} showing an overabundance of refractory elements \\citep[see a discussion on the solar system by][]{lod2003}, and (2) the direct spectroscopic signature of gaseous discs as in the cases of SDSS~J122859.93+104032.9 and SDSS~J104341.53+085558.2 \\citep{gan2006,gan2007} or dusty discs \\citep[e.g.,][]{far2009,bri2009}. The accreted material may also be supplied by mass loss from a close, low-mass, and possibly substellar companion. Only a few substellar companion to white dwarf stars are known such as the white dwarf plus L8-9 resolved pair PHL~5028 \\citep{ste2009} or close DA plus L8 post-common envelope binary WD0137$-$349 \\citep{max2006,bur2006}. Relatively luminous DAZ white dwarfs such as EG~102 \\citep{hol1997} may hide low-mass companions, but \\cite{deb2005} limits them to substellar types. The mass loss material captured by the deep potential well of a white dwarf should reflect the composition of the chromosphere of the red dwarf, which, in most cases, should be close to solar. In this context, the discovery of the peculiar DAZ white dwarf GALEX~J1931+0117 is timely. We present a series of ultraviolet (UV), optical and near infrared (NIR) observations (Section 2) and our model atmosphere analysis (Section 3.1) revealing a DAZ white dwarf with a peculiar abundance pattern (Section 3.2). A notable NIR excess (Section 3.3) may hold clues concerning the nature of this external source. We summarise and conclude in Section 4. ", "conclusions": "We identified and analysed the properties of one of the most heavily polluted white dwarfs known. The source of the material remains unknown, although the measured NIR excess may be attributed to a L5 dwarf or to a warm debris disc. The measured abundances suggest that the material is accreted onto the atmosphere in solar proportions and favour a model involving a L dwarf companion in a close orbit. However, the H$\\alpha$ emission notable in such systems is absent in GALEX~J1931+0117 and the mass loss imposed on the companion appears excessive. NIR intermediate dispersion spectroscopy should help determine the nature of the accretion source. Additional optical spectra will be useful to trace putative orbital motions while Space Telescope Imaging Spectrograph high-dispersion spectra will be useful to constrain further the carbon abundance and extend the pattern to less abundant elements." }, "1002/1002.0293_arXiv.txt": { "abstract": "The environment near the massive black hole (MBH) in the Galactic center is very hostile for star formation. Nevertheless, many young stars (both O and B stars) are observed close the MBH. The B-stars seems to have an isotropic, continuous distribution between 0.01 pc and up to a pc. The O stars, in contrast, seem to be distributed in a coherent disk like configuration, extending only between $\\sim0.04$ pc to $\\sim0.5$ pc. Our current understanding favors an in-situ formation origin for the more massive (O and Wolf-Rayet) stars, in gaseous disk and/or streams from an infalling gas clump. The B-stars seem to have a different origin, more likely through a dynamical capture, following binary disruption by the MBH. This scenario could also be able to explain the origin of hypervelocity stars in the Galactic halo. These and other possible origins of the young stars in the Galactic center are briefly reviewed and their possible observational signatures and constraints are detailed. ", "introduction": "High resolution observations have revealed the existence of many apparently normal young OB (including Wolf-Rayet; WR) stars in the galactic center (GC), where tidal forces exerted by the massive black hole (MBH; \\citep{sch+02b,ghe+03a}) are likely to inhibit regular star formation in regular molecular clouds. The existence and properties of such stars could give us important clues for understanding of the GC environment \\citep[see][for a review]{ale05}. They could also help constrain the origin and evolution of hypervelocity stars observed in the Galactic outskirts \\citep{hil88,bro+06a,per09}, which in principle could probe the potential and dark matter component of the Galaxy \\citep{gne+05,yuq+07,per+09}. The young stars observed in the central pc near the MBH, could be divided into two seemingly distinct stellar populations, which differ both in their types (B-stars vs. O or WR stars) and in their kinematic properties. The young B-stars population ($\\sim7-15\\, M_{\\odot}$; fainter stars can not currently be resolved) includes a few tens of stars with an isotropic distribution extending from $\\sim0.01$ pc all through the central pc \\citep{bar+10}. For some of these stars, in the central 0.04 pc (so called the 'S-stars' or the 'S-cluster'), full orbital solutions are known, showing them to have relatively high eccentricities ($0.3\\le e\\le0.95$) \\citep{ghe+03a,eis+05}, with approximately thermal eccentricity distribution, and random orbital orientations. Although the full kinematic properties of the B-stars outside this region are not known, the available data suggest their distribution is isotropic with similarly relaxed eccentricity distribution. The other young stars (mostly O and WR stars) reside in a more restricted region, between 0.04 pc to 0.5 pc, in seemingly two coherent structures. Most of these reside in a stellar disk moving clockwise in the gravitational potential of the MBH \\citep{lev+03,gen+03a,lu+06,pau+06,tan+06,bar+09}. The second structure is less coherent and its exact nature (and existence) is still unknown and debated. The orbits of the stars in the CW disk have an average eccentricity of $\\sim0.35$ and the opening of the disk is $h/R\\sim0.1$, where $h$ is the disk height and $R$ is its radius. The disk structure is warped at large angles ($65^{\\circ}$). Most of the stars outside the CW disk reside in somewhat less coherent structure, between 0.3-0.5 pc from the MBH, and highly inclined with respect to the CW disk. A small fraction of the young O stars do not reside in either of these structures at some intermediate inclinations. The current knowledge on the observed properties of the young stars in the GC are discussed in several recent papers (see e.g. \\citealp{bar+10}). In table 1, we summarize the observed properties of the young stars in the GC, which should be satisfied by models of their origin and evolution. {\\scriptsize }% \\begin{table} {\\scriptsize }\\begin{tabular}{|l|l|l|l|l|} \\hline {\\tiny Stellar} & {\\tiny Masses and } & {\\tiny Morphology} & {\\tiny Radial ($n(r)$)} & {\\tiny Eccentricity}\\tabularnewline {\\tiny Type} & {\\tiny Lifetime/age} & & {\\tiny Distribution} & {\\tiny Distribution}\\tabularnewline \\hline \\hline {\\tiny B} & {\\tiny 7-15 $M_{\\odot}$ } & {\\tiny Isotropic} & {\\tiny $r^{-1.1}$} & {\\tiny $\\sim$Thermal}\\tabularnewline & {\\tiny 10-50 Myrs} & {\\tiny (spherical)} & & {\\tiny $\\left\\langle e\\right\\rangle \\sim0.7$}\\tabularnewline & {\\tiny regular IMF} & & & \\tabularnewline \\hline {\\tiny O and WR} & {\\tiny 15-60 $M_{\\odot}$;} & {\\tiny Clockwise warped ($65^{\\circ}$) disk } & {\\tiny $r^{-2}$} & {\\tiny $\\left\\langle e\\right\\rangle \\sim0.35$}\\tabularnewline & {\\tiny{} 4-6 Myrs} & {\\tiny ($H/R\\sim0.1$) + coherent highly } & {\\tiny (in the CW disk)} & {\\tiny (in the CW disk)}\\tabularnewline & {\\tiny top heavy IMF} & {\\tiny inclined disk/stream structure} & & \\tabularnewline \\hline \\end{tabular}{\\scriptsize \\par} {\\scriptsize \\caption{{\\scriptsize Properties of the GC young stars}} } \\end{table} {\\scriptsize \\par} In the following we briefly overview suggested scenarios for the origin of the young stars in the Galactic center. Several models suggested these stars to be older stars, which only appear to be young. However, current observations show the young GC stars are apparently normal, genuinely young, and massive stars. We therefore focus on other more favorable scenarios, producing \\emph{young} stars (see \\citet{ale05} for an overview of the young stars impostors scenario, as well as some of less recent literature on some of the other models). ", "conclusions": "In these proceedings we have shortly reviewed the origin and evolution of the young stars in the Galactic center. These stars which could be divided into two distinct stellar populations likely originated from two different processes. The young stellar disk which contains mostly O and WR stars likely originated from an in-situ star formation through fragmentation of an infalling gaseous clamps. The population of young B-stars isotropically distributed throughout the central pc around the MBH likely have a different origin. Such stars were not likely to have been produced like the O-stars, and then migrated to their current postions, since a much larger parent population of B-stars should have been observed in the central pc. A binary disruption scenario, in which binaries which formed outside the central pc were scattered onto the MBH, could still be consistent with current observations. Such a scenario have specific predictions regarding the kinematic properties of the GC B-stars. We have reviewed these predictions which could be tested through direct observations in the coming few years. \\vspace{-0.2cm}" }, "1002/1002.2296_arXiv.txt": { "abstract": "We report the detection of an extended X-ray nebulosity with an elongation from northeast to southwest in excess of 15$^{\\prime\\prime}$ in a radial profile and imaging of the recurrent nova T Pyx using the archival data obtained with the X-ray Multi-Mirror Mission (\\XMM), European Photon Imaging Camera (pn instrument). The signal to noise ratio (S/N) in the extended emission (above the point source and the background) is 5.2 over the 0.3-9.0 keV energy range and 4.9 over the 0.3-1.5 keV energy range. We calculate an absorbed X-ray flux of 2.3$\\times$10$^{-14}$ erg cm$^{-2}$ s$^{-1}$ with a luminosity of 6.0$\\times$10$^{32}$ erg s$^{-1}$ from the remnant nova in the 0.3-10.0 keV band. The source spectrum is not physically consistent with a blackbody emission model as a single model or a part of a two-component model fitted to the \\XMM data ({$kT{\\rm _{BB}}$} $>$ 1 keV). The spectrum is best described by two MEKAL plasma emission models with temperatures at 0.2$^{+0.7}_{-0.1}$ keV and 1.3$^{+1.0}_{-0.4}$ keV. The neutral hydrogen column density derived from the fits is significantly more in the hotter X-ray component than the cooler one which we may be attributed to the elemental enhancement of nitrogen and oxygen in the cold material within the remnant. The shock speed calculated from the softer X-ray component of the spectrum is 300-800 km s$^{-1}$ and is consistent with the expansion speeds of the nova remnant derived from the $Hubble\\ Space\\ Telescope$ (\\HST) and ground-based optical wavelength data. Our results suggest that the detected X-ray emission may be dominated by shock-heated gas from the nova remnant. ", "introduction": "Classical novae (CNe) outbursts are the explosive ignition of accreted material on the surface of the white dwarf (WD) in a cataclysmic variable (CV) as a result of a thermonuclear runaway causing the ejection of 10$^{-7}$ to 10$^{-3}$ M$_{\\odot}$ of material at velocities up to several thousand kilometers per second (Shara 1989; Livio 1994; Starrfield 2001; Bode $\\&$ Evans 2008). Though there has only been one previous (and one very marginal) detection of old CNe remnants in the X-ray wavelengths (Balman \\& \\\"Ogelman 1999; Balman 2005,2006; Pek\\\"on \\& Balman 2008), CNe have been detected in the hard X-rays (above 1 keV) as a result of accretion, wind-wind and/or blast wave interaction {\\it in the outburst stage} (O'Brien et al. 1994; Krautter et al. 1996; Balman, Krautter, \\\"Ogelman 1998, Mukai \\& Ishida 2001, Orio et al. 2001; Ness et al. 2003; Hernanz \\& Sala 2002,2007; Page et al. 2009). Recurrent novae (RNe) are a type of CNe where outbursts recur with intervals of several decades (Webbink et al. 1987; Hachisu $\\&$ Kato 2001; Bode \\& Evans 2008). In general, these systems are expected to have high accretion rates of 10$^{-8}$ to 10$^{-7}$ M$_{\\odot}$ yr$^{-1}$ onto massive WDs close to the Chandrasekhar limit; occurrence of recurrent outbursts in relatively less massive WDs is also possible (Starrfield, Sparks, Truran 1985; Prialnik \\& Kovetz 1995; Yaron et al. 2005). RNe are detected in the hard X-rays {\\it in the outburst stage} as a result of wind-wind and/or blast wave interaction during the outburst stage (Orio et al. 2005; Greiner \\& Di Stefano 2002; Bode et al. 2006; Sokoloski et al. 2006; Drake et al. 2009; Ness et al. 2009). Recently, extended X-ray emission associated with the radio jet of the recurrent nova RS Oph is discovered (Luna et al. 2009). Recurrent nova T Pyx has 5 recorded outbursts in 1890, 1902, 1920, 1944, and 1966 (Webbink et al. 1987). Ground-based CCD observations (Shara et al. 1989) show the existence of at least two of the shells extending to a size of $\\sim$ 20$^{\\prime\\prime}$ in diameter and also a faint [OIII] shell has been found. $Hubble\\ Space\\ Telescope$ (\\HST; 1994-2007) and ground-based imaging of the shell around T Pyx show thousands of knots in H$\\alpha$ and [NII] with expansion velocities of about 350-715 km s$^{-1}$, with shell expansion speed around 500 km s$^{-1}$ (Shara et al. 1997; O'Brien \\& Cohen 1998; Schaefer, Pagnotta \\& Shara 2010). The \\HST\\ observations support an interacting shells model producing the clumping, shock heating, and emission lines. Schaefer et al. (2010) show that most of the knots have not decelerated and are powered by the RN outbursts and originate from a CN outburst of the year 1866. T Pyx is suggested to be a wind-driven source (due to the high mass accretion rate of $\\dot M$ $\\sim$ 1$\\times$ $10^{-7}$ M$_{\\odot}$ yr$^{-1}$) and a Super Soft X-ray source (SSS) (Patterson et al. 1998). On the other hand, Greiner \\& Di Stefano (2002) reports a ROSAT non-detection of the source in December 1998 excluding the possibility of the existence of a SSS. Gilmozzi \\& Selvelli (2007) and Selvelli et al. (2008) show that the UV+opt+IR spectrum of T Pyx is dominated by the accretion disk and the continuum in the UV can be represented by a blackbody of T $\\sim$ 34000 K with $\\dot M$ $\\sim$ 1$\\times$10$^{-8}$ M$_{\\odot}$ yr$^{-1}$. Their detailed study based on the UV data excludes the possibility that T Pyx is a SSS and Selvelli et al. (2008) uses this X-ray Multi-Mirror Mission (\\XMM) data set to show that a SSS nature is not supported. ", "conclusions": "The spectrum of T Pyx is well described by two different plasmas in ionization equilibrium at different temperatures. We stress that the expected SSS is not detected and the radial profiles deviate significantly from the PSF of EPIC pn. If one assumes all the detected luminosity is due to accretion, this yields an accretion rate less than $<$ a few$\\times$10$^{-10}$ M$_{\\odot}$ yr$^{-1}$ for the system in contradiction with the optical and UV measurements by a factor of 10-100. This strengthens the possibility that most of the detected emission is of the nova remnant in the X-rays rather than the point source. A simple shocked-shell model is of thermal origin. The total power from the shocked-shell, as an X-ray emitting nebula, can be expressed as in Balman (2005, 2006) $ L_x\\simeq 3.1\\times 10^{33} T_{7}^{0.5} n_o^2 R_{18.5}^{3}$. The temperature $T$ is in units of 10$^7$ K and radius of the shell $R$ is in units of 3.1$\\times 10^{18}$ cm. Using $R\\sim4.2\\times 10^{17}$ cm, $n_o\\sim$1-50 cm$^{-3}$ and $T\\sim 10^{6.3-7.3}$ K, $L_{\\rm x}$ is about 1.0$\\times 10^{31}$-8.0$\\times 10^{33}$ ergs s$^{-1}$. The detected X-ray luminosity (with \\XMM) is in this expected range of $L_{\\rm x}$ for the shocked-shell emission. The detected emission measure $EM = ^2 V_{eff}$ (calculated from normalisation of the fit) yields an average electron density $n_e$ of about 4.7 cm$^{-3}$ and 5.5 cm$^{-3}$ for the colder and hotter plasma, respectively, using a volume of 3.0$\\times 10^{53}$ cm$^{3}$ (consistent with a spherical region of 8$^{\\prime\\prime}$ radius at 3.5 kpc) and a filling factor f=1. If the filling factor is as low as 1$\\times 10^{-4}$, then the electron density can be as high as 470-550 cm$^{-3}$ in the X-ray emitting region. The electron density calculated for the nova shell of GK Per is of similar order 0.6-11.0 cm$^{-3}$ for f=1 (Balman 2005). The spectrum of the nova remnant of GK Per also shows two plasma emission components with a significant difference of absorption between the two components (Balman 2005). Such differences can be explained by enhanced non-solar abundances of metals in cold material (possibly a shell or collection of dense knots) between the two emission components. A simple fit with the VARABS or VPHABS model (within XSPEC) using variable abundances for the second harder X-ray component yields an $N{\\rm _{H2}}$ that is similar to $N{\\rm _{H1}}$ within error ranges with enhanced abundances of N/N$_{\\odot}$$\\sim$ 20 and O/O$_{\\odot}$$\\sim$ 10 . The remnant of T Pyx shows strong [NII] emission, but the abundances of metals is unknown. The $N{\\rm _{H}}$ difference can also be attributed to some complicated warm absorber effect (e.g. on the central binary system or the inner shell). Contini \\& Prialnik (1997) have modeled the circumstellar interaction of T Pyx shells. They find that as the latest shell catches the older shell, a forward shock moves into the older ejecta and a reverse shock moves into the new ejecta with typical electron densities of 100-300 cm$^{-3}$ and electron temperatures of about 0.1-0.9 keV. This model is very consistent with the X-ray spectral parameters obtained from the EPIC pn data (soft component). The plasma temperatures of the two different MEKAL model components can be used to calculate the shock velocities using the general relation $kT_s=(3/16)\\mu m_H (v_s)^2$, ($T = 1.4\\times10^5 v^2_{\\ 100 km s^{-1}}$), assuming Rankine-Huguenot jump conditions in the absence of particle acceleration. We derive 400 km s$^{-1}$ (300-800 km s$^{-1}$ using error ranges) for the first plasma emission component and 1050 km s$^{-1}$ for the second more absorbed plasma emission component (maximum limit 1400 km s$^{-1}$ using error ranges). The shell expansion velocity is measured to be about 350-715 km s$^{-1}$ (Shara et al. 1989; Shara et al. 1997;O'Brien \\& Cohen 1998; Schaefer et al. 2010). This is consistent with the shock speeds calculated using the spectral parameters of the softer X-ray component. The radial profiles calculated from the \\HST\\ data require a multiple shell model with a particular shell found around 5\\arcsec\\-6\\arcsec\\ and an extended emission region that goes out to 10\\arcsec (both measured from the position of T Pyx; Shara et al. 1997; Schaefer et al. 2010). The [NII]+H${\\alpha}$ images evidently shows an elongation/extension from the northeast to the southwest of the source position (see Figure 2c in Shara et al. 1989). The X-ray image indicates a similar elongation from the northeast to the southwest as well. This may be a result of the interaction between a bipolar outflow (e.g., a fast wind) during a RN outburst and the circumstellar environment (e.g., different shells) of the nova or the suggested CN remnant of 1866 by Schaefer et al. (2010). The expansion speeds of the 1966 outburst are in a range 850-2000 km s$^{-1}$ (Catchpole 1969). It has been about 40 years since the last outburst and the more absorbed (i.e. embeded) plasma emission component may belong to the most recent outburst. The expected location of the ejecta is about 2\\arcsec\\ - 4\\arcsec\\ (from the position of T Pyx). The expected size of the new interaction zone is within the core size of the EPIC pn PSF. The harder X-ray component may, also, belong to the binary system. A long observation of this source using the $Chandra$ Observatory (yielding better statistics) can resolve this issue with the aid of its superb pixel and PSF resolution." }, "1002/1002.3365_arXiv.txt": { "abstract": "We use the multi-epoch, mid-infrared {\\it Spitzer} Deep, Wide-Field Survey to investigate the variability of objects in 8.1 deg$^2$ of the NOAO Deep Wide Field Survey Bo{\\\"o}tes field. We perform a Difference Image Analysis of the four available epochs between 2004 and 2008, focusing on the deeper 3.6 and 4.5 \\micron~bands. Out of $474179$ analyzed sources, 1.1\\% meet our standard variability selection criteria, that the two light curves are strongly correlated ($r>0.8$) and that their joint variance ($\\sigma_{12}$) exceeds that for all sources with the same magnitude by $2 \\sigma$. We then examine the mid-IR colors of the variable sources and match them with X-ray sources from the XBo{\\\"o}tes survey, radio catalogs, 24 \\micron\\, selected active galactic nucleus (AGN) candidates, and spectroscopically identified AGNs from the AGN and Galaxy Evolution Survey (AGES). Based on their mid-IR colors, most of the variable sources are AGNs (76\\%), with smaller contributions from stars (11\\%), galaxies (6\\%), and unclassified objects, although most of the stellar, galaxy, and unclassified sources are false positives. For our standard selection criteria, 11\\%--12\\% of the mid-IR counterparts to X-ray sources, 24 \\micron~AGN candidates, and spectroscopically identified AGNs show variability. The exact fractions depend on both the search depth and the selection criteria. For example, 12\\% of the 1131 known $z>1$ AGNs in the field and 14\\%--17\\% of the known AGNs with well-measured fluxes in all four IRAC bands meet our standard selection criteria. The mid-IR AGN variability can be well described by a single power-law structure function with an index of $\\gamma\\approx0.5$ at both 3.6 and 4.5 \\micron, and an amplitude of $S_0\\simeq0.1$ mag on rest-frame timescales of 2 years. The variability amplitude is higher for shorter rest-frame wavelengths and lower luminosities. ", "introduction": "\\label{sec:intro} A substantial fraction of objects in the universe change in brightness with time. This apparent variability can be due to physical periodic changes in the objects (e.g., Cepheids, RR Lyrae), motion (e.g., eclipsing binaries, microlensing, rotating stars), explosive events (e.g., supernovae, novae), and accretion (e.g., cataclysmic variables, active galactic nuclei (AGNs)). In a high-latitude extragalactic survey, the most common variable sources are active galactic nuclei (e.g., \\citealt{2007AJ....134.2236S}). Optical variability is one method for identifying quasars\\footnote{We will use words ``quasar'' and ``AGN'' interchangeably throughout this paper.} (e.g., \\citealt{2002AcA....52..241E,2003AJ....125....1G,2004ApJ...606..741R}), although it is only recently that a fully quantitative approach to variability selection has been developed (\\citealt{2010ApJ...708..927K}). The emission from quasars generally has three components. The UV to near-IR radiation is dominated by a hot accretion disk extending from an inner edge of a few gravitational radii from the black hole outward with, in simple thin disk theory, a temperature profile $T \\propto R^{-3/4}$ (\\citealt{1973A&A....24..337S}). Near the inner edge, there is a corona of hotter gas that produces the non-thermal X-ray emission (\\citealt{1994ApJ...432L..95H}). On scales where the temperature is below the dust sublimation temperature ($\\sim 2000$~K), dust absorbs radiation from the disk and reradiates the energy in the mid-IR and far-IR (\\citealt{1987ApJ...320..537B}). The overall spectrum typically shows a minimum near 1 \\micron, with the emission from the disk rising toward the UV and the emission reradiated by dust rising toward the far-IR (e.g., \\citealt{1989ApJ...347...29S}). There is an increasing evidence that many physical relations between AGNs, galaxies and their large scale clustering have to be taken into account if their formation and evolution is to be understood (e.g., \\citealt{2008A&A...490..893T}). We know a great deal about the optical variability of quasars both from large studies of the variability seen in ensembles of sparsely monitored quasars and from detailed studies of individual quasars. Ensemble studies (e.g., \\citealt{2004ApJ...601..692V,2005AJ....129..615D}) have shown that variability increases with decreasing optical wavelength, decreasing luminosity, and potentially decreasing black-hole mass. The structure function of the ensemble variability is a power law with smaller variability amplitudes on short timescales, with some evidence for saturation on timescales of order a few decades. Until recently, there were few studies of individual quasars (e.g., \\citealt{1985ApJ...296..423C,1989ApJ...337..236C,1994MNRAS.268..305H}), but this has changed dramatically in the last year. The light curves of individual quasars are well modeled by a damped random walk, a stochastic process described by the amplitude of the random walk and a damping timescale for returning to the mean luminosity (\\citealt{2009ApJ...698..895K,2010ApJ...708..927K,2010arXiv1004.0276M}). While preliminary indications from \\cite{2009ApJ...698..895K}, based on $\\sim100$ quasars, suggest that these two process parameters are related to the quasar luminosity and black hole mass, \\cite{2010arXiv1004.0276M} used $\\sim9000$ Sloan Digital Sky Survey (SDSS) Stripe 82 quasars to find a number of clear trends. For example, the asymptotic variability on long timescales decreases with increasing luminosity and rest-frame wavelength, and is correlated with black hole mass. The timescale for returning to the mean luminosity increases with wavelength and also with increasing black hole mass, but remains constant with redshift and luminosity. Far less is known about the near-IR and mid-IR variability of quasars. Since AGN luminosities do vary in time at optical wavelengths, we expect to find some variability in the IR, but it could be heavily smoothed by averaging the response over the large scales of either the cooler parts of the disk or the absorbing dust. \\cite{2004MNRAS.350.1049G} found that the majority (39/41) of the surveyed low-luminosity nearby Seyferts do vary in the near-IR and mid-IR, and that the variability is most apparent at longer wavelengths due to diminishing flux from the host galaxy. The most extensive work to date on near-IR and mid-IR variability of quasars is that of \\cite{1999AJ....118...35N}. These authors studied 25 low-redshift quasars ($z\\simeq0.1$), which were observed in five bands (1.3 -- 10.6 \\micron) for $\\sim30$ years. The amplitude of variations for radio loud (quiet) quasars is $\\sim0.3$ mag (0.1 mag) on rest-frame time lags (time differences between any two epochs) of 2 years at 10.6 \\micron. For one quasar, PG 1226+023, they measured structure functions in five bands, all showing $\\sim0.1$ mag variations for a rest-frame time lag of 2 years. Another example is the work of \\cite{2006ApJ...639...46S}, who found a delayed response between the $K$- and $V$-band variability in four Seyfert 1 galaxies. These authors showed that $K$-band light curves have a fairly tight time lag--luminosity relation, $\\tau \\approx L^{1/2}$, while broad line region (BLR) time lags are more scattered. The interpretation is that the $K$-band emission is mainly coming from dusty clouds bounding the BLR, and that the size scale of these dust clouds is very dependent on the AGN luminosity, since they have to be colder than the dust sublimation temperature. In this paper, for the first time, we investigate the variability of a truly large number of mid-IR sources, using the four-epoch {\\it Spitzer} Deep Wide Field Survey (SDWFS; \\citealt{2009ApJ...701..428A}) of the Bo{\\\"o}tes field of the NOAO Deep Wide Field Survey (NDWFS; \\citealt{1999ASPC..191..111J}). The two goals are to characterize the mid-IR variability of a large sample of AGNs, and to use the variability to identify additional AGNs. While mid-IR quasar selection (\\citealt{2005ApJ...631..163S}) works very well when the AGN dominates the luminosity (even for quasars behind dense stellar fields such as LMC/SMC; \\citealt{2009ApJ...701..508K}), the method begins to fail as the luminosity becomes comparable to that of the host (\\citealt{2008ApJ...679.1040G,2010ApJ...713..970A}). With mid-IR variability, we should be able to identify such sources and extend the sample of mid-IR selected quasars to lower luminosities. We also make use of the extensive wavelength coverage of the field, including X-ray (XBo\\\"otes; \\citealt{2005ApJS..161....1M}), the earlier Infrared Array Camera (IRAC) Shallow Survey that became the first epoch of SDWFS \\citep{2004ApJS..154..48E}, the 24 \\micron~survey of the field (see \\citealt{2005ApJ...622L.105H}), and the FIRST \\citep{1995ApJ...450..559B} and WSRT \\citep{2002AJ....123.1784D} radio surveys. Based on these data, the AGN and Galaxy Evolution Survey (AGES; C. S. Kochanek et al. 2010, in preparation) selected galaxies and AGNs at all these wavelengths for spectroscopic follow up using the Hectospec instrument on the MMT (\\citealt{2005PASP..117.1411F}), providing roughly 23,500 redshifts in the field. In Section~\\ref{sec:data} we describe the data analysis, and in Section~\\ref{sec:definitions} we describe our variability selection criteria. In Section~\\ref{sec:results}, we present the results of this study. The mid-IR structure function analysis is presented in Section~\\ref{sec:sfunction} followed by a summary in Section~\\ref{sec:summary}. Throughout this paper, we use a standard $\\Lambda$CDM model with $(\\Omega_{\\Lambda}, \\Omega_{\\rm M}, \\Omega_{k}) = (0.7, 0.3, 0.0)$ and $h=H_0/100=0.70$. Where needed, we use the galaxy and AGN templates of \\cite{2010ApJ...713..970A} to describe the mid-IR colors of various populations and to make any {\\it K}-corrections or estimates of absolute luminosities. ", "conclusions": "\\label{sec:summary} In this paper, we performed DIA photometry on the four epochs of SDWFS data. The common area of the {\\it Spitzer} mosaics covers 8.1 deg$^2$ of the NDWFS Bo{\\\"o}tes field, and contains $474,179$ mostly extragalactic objects. We concentrate on the deeper 3.6 and 4.5 \\micron~channels, and provide variability catalogs and light curves. These variability data are cross-correlated with the AGES redshift survey, X-ray and radio data, and 24 \\micron~AGN candidates from the same area of the sky to study the variability of various classes of objects. We study the mid-IR variability of objects based on the significance level of their joint variability ($\\sigma_{12}$) in both the [3.6] and [4.5] channels focusing on the sources that also show a high correlation ($r>0.8$) between the two channels. Our standard selection criteria with $\\sigma_{12}>2$ identify $\\sim5100$ objects as variables, which constitute 1.1\\% of all detected sources in the field. We find that the majority (76\\%) of the mid-IR variable sources in the NDWFS Bo{\\\"o}tes field are AGNs, where we define AGNs to be objects that satisfy any of the following criteria: (1) lie inside the modified AGN wedge, (2) are X-ray, 24 \\micron~or radio AGN candidates, (3) have $z>1$, or (4) have an AGN spectroscopic template in any region. Amongst all the extragalactic classes of objects, AGNs are the primary sources able to produce significant changes in their mid-IR luminosity over a several year timescale. Thus, mid-IR variability may be used as a new tool for selecting AGNs, especially if combined with other methods. However, in our relatively shallow few-epoch survey, only $\\sim$ 15\\% of AGNs are sufficiently variable to be detected; therefore, variability-selected samples of AGNs are highly incomplete. This incompleteness is caused by the typical variability amplitude being comparable to the survey photometric errors, and is not due to the small number of available epochs. For our standard criteria, 14\\% of X-ray sources, 17\\% of 24 \\micron\\ selected AGN candidates, and 15\\% of spectroscopically confirmed AGNs are detected by their variability. The variable AGNs primarily occupy the mid-IR AGN selection wedge of \\cite{2005ApJ...631..163S}, with an extension to bluer $[3.6]-[4.5]$ colors where the host galaxy dominates the mid-IR colors of low-luminosity AGNs (see \\citealt{2008ApJ...679.1040G,2010ApJ...713..970A}). If we redefine the original \\cite{2005ApJ...631..163S} selection wedge to encompass X-ray or variability-selected AGNs, we find that variable objects bluer than $[3.6]-[4.5]>0.3$ mag are likely to be AGNs. This opens a new window for the {\\it Spitzer} Warm Mission, where the $[3.6]-[4.5]>0.3$ mag color cut combined with the variability criterion will be a solid indication of AGN activity. The structure function of $z>1$ AGNs is well described by a power law with a logarithmic slope of $\\gamma = 0.56 \\pm 0.18$ and amplitude $S_0=0.11\\pm0.02$ mag in [3.6] and $\\gamma = 0.44 \\pm 0.14$, $S_0=0.12\\pm0.01$ mag in [4.5] for a rest-frame time lag of $\\tau_0=2$ years. These structure functions are steeper than those observed in optical bands, and we argue that this is due to the inability of either disk or dust emission to respond on the short timescales seen in optical studies. The mid-IR structure functions go approximately as $\\tau^{1/2}$, which corresponds to the short timescale ($\\tau \\ll \\tau_{\\rm model}$) behavior in the damped random walk model used by \\cite{2009ApJ...698..895K}, \\cite{2010ApJ...708..927K}, and \\cite{2010arXiv1004.0276M}. If we examine the variability of various subsamples, we detect more variability in low luminosity objects observed at shorter rest-frame wavelengths. The mid-IR sources corresponding to X-ray and 24 \\micron~AGN have structure functions identical to that of $z>1$ AGNs. However, this is not the case for the counterparts of radio sources which have $S_0\\simeq0.05\\pm0.01$ mag in both [3.6] and [4.5] bands, as compared to $S_0\\simeq0.10\\pm0.01$ for either $z>1$ objects, X-ray, or MIPS sources. One explanation may be that the radio sample has more contamination by non-AGN (e.g., due to confusion between radio cores and lobes), another is that the radio-selected AGNs are accreting in a different mode (lower accretion rates or efficiencies; e.g., \\citealt{2008A&A...490..893T}) and that variability in such a mode is inherently lower. Expanding the SDWFS survey with new epochs during the {\\it Spitzer} Warm Mission would represent a significant improvement. Given our structure functions, we can estimate the results from expanding the SDWFS survey with new epochs. We considered two cases of (1) doing one new epoch in 2011 March and (2) doing new epochs in both 2011 and 2012 March. The improvements in the mean $\\chi^2$ are significant, with $\\langle \\chi^2 \\rangle = 4 + 3.0 S_0^2/\\sigma^2$ and $5 + 4.3 S_0^2/\\sigma^2$, respectively, for AGNs, compared to $\\langle \\chi^2 \\rangle = 4$ or 5 for non-variable objects. These represent a significant improvement over the SDWFS baseline, where $\\langle \\chi^2 \\rangle = 3 + 1.6 S_0^2/\\sigma^2$ for AGN and 3 for non-variable objects. Figure~\\ref{fig:compl} shows the increases in completeness for these five and six epoch scenarios. Most of the gain for detecting variable sources comes from adding one additional epoch. The greatest astrophysical gain from two additional epochs comes from being able to better map out and subdivide the structure function by luminosity, redshift, and source type to better characterize the physics of AGN variability at these wavelengths. In particular, it would confirm the different slope of the mid-IR structure function from that observed in the optical." }, "1002/1002.1360_arXiv.txt": { "abstract": "In this paper, results of 2.5-dimensional magnetohydrodynamical simulations are reported for the magnetic reconnection of non-perfectly antiparallel magnetic fields. The magnetic field has a component perpendicular to the computational plane, that is, guide field. The angle $\\theta$ between magnetic field lines in two half regions is a key parameter in our simulations whereas the initial distribution of the plasma is assumed to be simple; density and pressure are uniform except for the current sheet region. Alfv\\'en waves are generated at the reconnection point and propagate along the reconnected field line. The energy fluxes of the Alfv\\'en waves and magneto-acoustic waves (slow mode and fast mode) generated by the magnetic reconnection are measured. Each flux shows the similar time evolution independent of $\\theta$. The percentage of the energies (time integral of energy fluxes) carried by the Alfv\\'en waves and magneto-acoustic waves to the released magnetic energy are calculated. The Alfv\\'en waves carry 38.9\\%, 36.0\\%, and 29.5\\% of the released magnetic energy at the maximum ($\\theta=80^\\circ$) in the case of $\\beta=0.1$, $1$, and $20$ respectively, where $\\beta$ is the plasma $\\beta$ (the ratio of gas pressure to magnetic pressure). The magneto-acoustic waves carry 16.2\\% ($\\theta=70^\\circ$), 25.9\\% ($\\theta=60^\\circ$), and 75.0\\% ($\\theta=180^\\circ$) of the energy at the maximum. Implications of these results for solar coronal heating and acceleration of high-speed solar wind are discussed. ", "introduction": "The heating mechanism of the solar corona is one of the most mysterious issue in astrophysics (Aschwanden 2004). The structure of the solar atmosphere consists of the photosphere, the temperature of which is about or more than six thousand degrees, the chromosphere, the temperature of which is up to a few ten thousand degrees, the transition region, in which the temperature increases rapidly, and the corona, the temperature of which reaches a few million degrees. To maintain such a high temperature in the corona in spite of cooling by heat conduction and radiative losses, a continuous supply of thermal energy is necessary. X-ray observations from early space experiments (e.g., Skylab) have shown that the corona is not uniform and consists of many bright loops. Poletto et al. (1975) demonstrated a correspondence between enhanced X-ray emission and a magnetic loop. It is therefore suggested that magnetic activity is related to the heating of the solar corona. Related to the magnetic activity, there are two promising models for coronal heating (see reviews by Klimchuk 2006, Erd\\'elyi \\& Ballai 2007 and also references therein). One is heating by the dissipation of Alfv\\'en waves \\footnote{Other types of waves (acoustic, slow-mode and fast-mode waves) are strongly damped or reflected at the steep density and temperature gradients of chromosphere and transition region.} that propagate in the magnetic flux tubes (Alfv\\'en 1947; Hollweg 1973, 1981, 1984, 1986; Uchida \\& Kaburaki 1974; Wentzel 1974; McKenzie, Banaszkiewicz, \\& Axford 1995; Axford et al. 1999). There are many proposed mechanisms for the dissipation of Alfv\\'en waves, such as mode conversion (see the next paragraph), resonant absorption (e.g., Ionson 1978; Poedts et al. 1989; Erd\\'elyi \\& Goossens 1995; Ruderman et al. 1997), phase mixing (e.g, Heyvaerts \\& Priest 1983), or magnetohydrodynamic (MHD) turbulence (e.g., Inverarity \\& Priest 1995; Matthaeus et al. 1999). The other is heating by many small-scale flares, i.e., nanoflares triggered by magnetic reconnection (Parker 1988). Observationally the occurrence frequency of microflares and nanoflares, $N$, has been found to be $dN/dW \\propto W^{-\\alpha}$, where $W$ is the total flare energy. The power-law index, $\\alpha$, ranges from 1.5 to 1.8 (Hudson 1991; Shimizu 1995; Shimojo \\& Shibata 1999; Aschwanden \\& Parnell 2002) and nanoflares can not account for coronal heating if $\\alpha$ is less than 2. However, a value of $\\alpha$ larger than 2 has also been found (Krucker \\& Benz 1998; Parnell \\& Jupp 2000). The definitive conclusion has not yet been obtained from observations. As an origin of Alfv\\'en waves, Kudoh \\& Shibata (1999) considered a photospheric random motion propagating along an open magnetic flux tube in the solar atomosphere, and performed 1.5-dimensional (1.5D, i.e., torsional motion is allowed) MHD simulations for solar spicule formation and the heating of the corona. It was shown that Alfv\\'en waves transport sufficient energy flux into the corona to account for its heating, by extending the work by Hollweg, Jackson, \\& Galloway (1982) (see also Saito, Kudoh, \\& Shibata 2001). Moriyasu et al. (2004) performed 1.5D MHD simulations of the propagation of nonlinear Alfv\\'en waves along a closed magnetic loop including heat conduction and radiative cooling. They found that the corona is episodically heated by fast- and slow-mode MHD shocks generated by nonlinear Alfv\\'en waves via nonlinear mode-coupling. It was also found that the time variation of the simulated extreme-ultraviolet (EUV) and X-ray intensities is quite similar to the observed one. They concluded that the observed nanoflares may not be a result of reconnection but may be due to nonlinear Alfv\\'en waves. Subsequently, Antolin et al. (2008) discussed the observational signatures of the power-law indexes and coronal heating mechanisms (Alfv\\'en waves and nanoflares) by using 1.5D MHD model of Moriyasu et al. (2004). They found that Alfv\\'en heating and nanoflare heating exhibit different power-law indexes (see also Antolin \\& Shibata 2010). However, since Alfv\\'en waves can be generated by magnetic reconnection unless the reconnection takes place in perfectly antiparallel magnetic fields, the nanoflare model may not be very different from the Alfv\\'en wave model. Yokoyama \\& Shibata (1995) modelled jets (X-ray or EUV jets and serges observed with H$\\alpha$ in the chromosphere) by performing a resistive two-dimensional MHD simulation of the magnetic reconnection occurring in the current sheet between emerging magnetic flux and overlying pre-existing coronal magnetic fields. Recent {\\it Hinode} (Kosugi et al. 2007) observation revealed that jets are ubiquitous in the chromosphere (Shibata et al. 2007). Nishizuka et al. (2008) proved that the jets on the basis of this model are quantitatively corresponding to the multiwavelength jets (see also Cirtain et al. 2007 and Liu et al. 2009) observed with {\\it Hinode} and {\\it TRACE} (Handy et al. 1999). Yokoyama (1998) found that the ratio of energy of Alfv\\'en waves to the energy released by the magnetic reconnection (the model of Yokoyama \\& Shibata 1995) is nearly equal to 3\\%. In this model, there is a shear between emerging magnetic fields and coronal magnetic fields, kinks are produced by the magnetic reconnection and propagate away as Alfv\\'en waves. Takeuchi \\& Shibata (2001) performed 2.5-dimensional (2.5D) MHD simulations of the photospheric magnetic reconnection caused by convection, and found that the energy flux of Alfv\\'en waves \\footnote{The perpendicular magnetic field injected by magnetic reconnection propagates as Alfv\\'en waves.} is enough to explain both coronal heating and spicule production. The generation of Alfv\\'en waves through magnetic reconnection has been also discussed by Parker (1991), Axford et al. (1999), and Sturrock (1999) in the context of coronal hole heating and solar wind acceleration. The upward acoustic waves in the flux tube are expected to be transported as other waves (e.g., slow-mode waves) after the formation of the shocks. Slow-mode waves are also thought to contribute to coronal heating though the contribution is only to the relatively low corona because of their compressibility. Takeuchi \\& Shibata (2001) measured the energy flux of slow-mode waves and found that it was ten times larger than that of Alfv\\'en waves. Suzuki (2002) discussed the possibility of coronal heating and the acceleration of the low-speed solar wind by slow-mode waves. Suzuki (2004) developed his study by including fast-mode waves (linearly polarized Alfv\\'en waves) and showed that slow-mode waves contribute to the heating of the low corona and the acceleration of the low-speed solar wind, while linearly polarized Alfv\\'en waves contribute to the heating of the outer corona and the acceleration of the high-speed solar wind. As the mention above, waves are created by torsional motions as recently observed (e.g., Bonet et al. 2008) in the photosphere or magnetic reconnections. They propagate into the corona and disipate their energy through linear and non-linear mechanisms. In particular, Alfv\\'en waves can be generated by such drivers, for the first time, by Jess et al. (2009). From {\\it Hinode} data, De Pontieu et al. (2007) estimated the energy flux carried by transversal oscillations generated by spicules and compared with radiative MHD simulations by more realistically extending Kudoh \\& Shibata (1999). They indicated that the calculated energy flux is enough to heat the quiet corona and to accelerate the high-speed solar wind. Okamoto et al. (2007) also estimated the energy flux to be $2\\times 10^6$ ergs cm$^{-2}$ s$^{-1}$ propagating on coronal magnetic fields. These reports (see also Tomczyk et al. 2007), however, have considerable argument about what is Alfv\\'en waves. Erd\\'elyi and Fedun (2007) and Van Doorsselaere et al. (2008) argued that these oscillations were likely to be kink oscillations from observed behavior (see also Goossens et al. 2009; Taroyan \\& Erd\\'elyi 2009). It is thought that the magnetic reconnection causes the solar flare. Particle acceleration takes place associated with solar flare. The mechanism of the particle acceleration, however, is not made clear. Alfv\\'en waves could contribute to the acceleration of ions through cyclotron resonance (see, e.g., Miller 2000, and references therein). The generation of Alfv\\'en waves by the magnetic reconnection is a very interesting research topic also from such a point of view. In this paper, we present the results of 2.5D MHD simulations of the Alfv\\'en wave generation by the reconnection of non-perfectly antiparallel magnetic fields. The initial magnetic fields have a shear, i.e., the magnetic field has a component perpendicular to the computational plane. This component can not contribute to the reconnection and is so-called guide field. The magnetic reconnection in this geometry was analytically studied by Petschek \\& Thorne (1967). The angle between the magnetic field lines in two half regions is a parameter, $\\theta$, while the initial distribution of the plasma is assumed to be simple. The energy fluxes of Alfv\\'en waves and magneto-acoustic waves are measured. In section 2 we describe the numerical method and model. In section 3 we show the results of the simulations. Finally a discussion and summary are given in section 4. ", "conclusions": "Figure \\ref{FIG08} shows the ratio (percentage) of the energies carried by the Alfv\\'en waves and magneto-acoustic waves ($E_\\mathrm{Alfven}$ and $E_\\mathrm{Sound}$) to the released magnetic energy at the final stage of simulation. In the case of $\\theta=180^\\circ$, $E_\\mathrm{Alfven}$ is equal to 0 independent of the plasma $\\beta$ because the magnetic field component perpendicular to the initial field is not generated by the reconnection. As $\\theta$ decreases the percentage of the energy carried by the Alfv\\'en waves increases up to 38.9\\%, 36.0\\%, and 29.5\\% in the case of $\\beta=0.1$, $1$, and $20$ respectively, and then decreases. The percentage is maximum when $\\theta = 80^\\circ$. The percentage of the energy carried by the magneto-acoustic waves is roughly constant when $\\theta$ is relatively large in the cases of $\\beta=0.1$ and $1$. In the case of $\\beta=20$, it gradually decreases in decreasing $\\theta$. The maximum values are 16.2\\%, 25.9\\%, and 75.0\\% in the case of $\\beta=0.1$, $1$, and $20$ respectively. While the percentage of the energy carried by the Alfv\\'en waves is almost independent of $\\beta$, that of the energy carried by the magneto-acoustic waves changes by some factor. The ratio of the energy carried by the magneto-acoustic waves to the released magnetic energy becomes larger as $\\beta$ becomes larger in $\\theta\\geq 40$. This means that the significant part of the energy released by the magnetic reconnection is transported as a perturbation of gas pressure in a high-$\\beta$ plasma case. This result can give a suggestion to the study of high-$\\beta$ plasma astrophysical objects, e.g., the accretion disk. The accretion disk is thought to be weakly magnetized. The magnetic reconnection is induced by the magnetorotational instability (e.g., Sano \\& Inutsuka 2001, Machida \\& Matsumoto 2003). Sano \\& Inutsuka (2001) showed that the heating rate is strongly related to the turbulent shear stress, which determines the efficiency of angular momentum transport. Therefore, the study of the magnetic reconnection in a high-$\\beta$ plasma is important. The total non-radiative energy input to the solar coronal hole was estimated at $5 \\times 10^5$ ergs cm$^{-2}$ s$^{-1}$ by Withbroe (1988). For the acceleration of high-speed solar wind, some $1 \\times 10^5$ ergs cm$^{-2}$ s$^{-1}$ is required to be deposited at distances of several solar radii (see, e.g., Parker 1991, and references therein). If the solar wind is accelerated by the energy flux of Alfv\\'en waves, this means that 20\\% of the energy released by reconnection events in the solar corona is transferred as a form of Alfv\\'en wave. Our results show that the energy larger than the required can be carried by the Alfv\\'en wave independent of $\\beta$ around the parametric region of $60^\\circ \\leq \\theta \\leq 110^\\circ$. These energies are converted to thermal energy through the dissipation of Alfv\\'en waves. The guide field, $B_z$ in this paper, can not contribute to the reconnection because of the 2.5D approximation. Besides this, the reconnection progresses typically with the Alfv\\'en time scale, which depends on the magnetic field on the $xy$-plane (more exactly speaking, depend on the reconnection component of the Alfv\\'en velocity). It is therefore interesting that how the results change when the normalizations are changed: Time is re-normalized by the effective Alfv\\'en time ($t_\\mathrm{A0,eff}=L_0/v_\\mathrm{A0,eff}$). This means that the same time in Figure \\ref{FIG06} is not the same time in Figure \\ref{FIG09} because the effective Alfv\\'en velocity is different according to $\\theta$. The magnetic field is also re-normalized by the initial magnetic field which can reconnect ($B_0 \\sin \\theta/2$). When $\\theta = 0$, the magnetic field is uniform and the reconnection does not take place. Figure \\ref{FIG09} shows the difference of the re-normalized magnetic energy from the initial value as the function of re-normalized time in each $\\theta$ case. In the high-$\\beta$ case, the magnetic energy is released at almost the same rate in the effective Alfv\\'en time when $\\theta \\ge 40^\\circ$. The lower $\\beta$ is, the larger the discrepancy of the release rate becomes. Figure \\ref{FIG10} shows the re-normalized energy fluxes and the time integral of those as the function of re-normalized time. The energy flux and its time integral of Alfv\\'en waves show the time variation similar to each other independent of $\\beta$. Those of magneto-acoustic waves also show the time variation similar to each other except when $\\theta$ is relatively small, although there is a bit of $\\beta$ dependence. Compared with Figure \\ref{FIG07}, the peak positions are matched. Figure \\ref{FIG11} indicates the percentage of the energies carried by the Alfv\\'en waves and magneto-acoustic waves to the released magnetic energy. The vertical axis is re-normalized as same as the above-mentioned way. This shows clearly that the amount of the energy carried by the Alfv\\'en waves has almost the same dependence on $\\theta$ independent of $\\beta$ {\\it if the energy and time are scaled by the effective magnetic energy and Alfv\\'en time}. This is equivalent to that the magnetic configuration is important rather than the field strength relative to the gas pressure for the energy release rate in the effective Alfv\\'en time. On the other hand, the amount of the energy carried by the magneto-acoustic waves shows the $\\beta$ dependence. In this paper, we have reported the results for 2.5D MHD simulations of the magnetic reconnection. When magnetic fields are non-perfectly antiparallel, the magnetic field component perpendicular to the initial field is generated and propagates as the Alfv\\'en wave. We have measured the energy fluxes of Alfv\\'en waves and magneto-acoustic waves. Magneto-acoustic waves are related to fast mode waves in the high-$\\beta$ case ($\\beta=20$) and slow mode waves in the low-$\\beta$ case ($\\beta=0.1$). The energy carried by the Alfv\\'en waves is more than 30\\% of the energy released by the magnetic reconnection at the maximum. That value satisfies the requirement for energy flux of Alfv\\'en waves necessary for acceleration of high-speed solar wind in the nanoflare coronal heating model. For more exact discussion, 3-dimensional simulations with more realistic plasma distribution are necessary. We would like to thank David H. Brooks and Takeru Suzuki for helpful comments. This study was initiated as a part of the ACT-JST summer school for numerical simulations in astrophysics and in space plasmas. Numerical computations were carried out on VPP5000 (project ID: yhk32b, rhk05b, and whk08b) and the general-purpose PC farm (project P.I. KT) at Center for Computational Astrophysics, CfCA, of National Astronomical Observatory of Japan. 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." }, "1002/1002.3153_arXiv.txt": { "abstract": "We present the first astrophysical measurement of the pressure of cold matter above nuclear saturation density, based on recently determined masses and radii of three neutron stars. The pressure at higher densities are below the predictions of equations of state that account only for nucleonic degrees of freedom, and thus present a challenge to the microscopic theory of neutron star matter. ", "introduction": " ", "conclusions": "" }, "1002/1002.0016_arXiv.txt": { "abstract": "We present color-magnitude and morphological analysis of 54 low-redshift ultraluminous infrared galaxies (ULIRGs; $0.01810^{12}\\rm L_{\\odot}$) were first discovered in a large number by {\\it Infrared Astronomical Satellite} (IRAS) two decades ago (Rowan-Robinson et al. 1984; Soifer et al. 1984). These objects are among the most luminous sources in the universe, with most of their energy radiated in the infrared. Although the powering sources of ULIRGs remain uncertain, there is evidence that both AGNs and dusty starbursts contribute to their high bolometric luminosities (see Sanders \\& Mirabel 1996; Lonsdale, Farrah, \\& Smith 2006 for reviews). The majority of local ULIRGs are mergers with tidal features commonly seen (e.g. Sanders et al. 1988; Melnick \\& Mirabel 1990; Murphy et al. 1996), and distant ULIRGs are the possible progenitors of today's massive ellipticals (e.g. Barnes \\& Hernquist 1996). Recent large optical photometric surveys show that the color distribution of galaxies are bimodal (Strateva et al. 2001; Hogg et al. 2003; Blanton et al. 2003c; Bell et al. 2004a,b). Color-magnitude diagram reveals a red sequence of early type galaxies with little star formation activity, and a blue cloud of late type galaxies that form stars. The characteristic number density in the standard luminosity function of Schechter (1976), ${\\phi_{*}}$, has doubled for the red sequence galaxies since $z\\sim1$, but hardly changes for the blue cloud, implying that the number and stellar mass of the red sequence galaxies have been gradually built up, whereas those of the blue cloud galaxies remain nearly constant (e.g., Bell et al. 2004b; Faber et al. 2007). One scenario of building up the most massive galaxies in the red sequence is major merging of disk galaxies (Bell et al. 2004b), and thus galaxies in the ``green valley'', the region between the red sequence and blue cloud in the color-magnitude diagram, may represent this transitional type. As ULIRGs are proposed to be disk-disk mergers that will finally evolve to massive ellipticals,% studying the color-magnitude relation of ULIRGs may shed light on the evolutionary scenarios of building up massive ellipticals. Galaxy morphologies indicate the distribution of baryonic matter in the galaxies, and are therefore tracers of its formation and evolution. The majority of local ULIRGs are merging systems in which disturbed morphologies such as tidal features are commonly observed. Recent deep sky surveys have generated enormous amounts of galaxy imaging data, and automated morphology measurement methods are required to analyze those large datasets with much higher efficiencies than visual inspections provide. There are two quantitative approaches to describing galaxy morphology: parametric methods and non-parametric methods. Parametric methods fit the light distribution of a galaxy to pre-defined formulae, such as Sersic index fitting (Blanton et al. 2003c) and bulge-to-disk ratio (Peng et al. 2002; Simard et al. 2002), and are usually inadequate to describe the morphologies of complex systems. Non-parametric methods, such as the concentration, asymmetry, and clumpiness (CAS) system (e.g. Isserstedt \\& Schindler 1986; Abraham et al. 1994; Wu et al. 1999; Conselice et al. 2000; Conselice 2003), do not rely on pre-defined functions, and may be the better methods for describing complex systems. Abraham et al. (2003) introduced the $Gini$ coefficient ($G$), which describes the relative intensity distributions of a galaxy. The $Gini$ coefficient is correlated with concentration and increases with the fraction of light in compact components. Lotz, Primack, \\& Madau (2004; hereafter LPM04) introduced $M_{20}$, which is the second moment of the brightest 20\\% of galaxy flux. $M_{20}$ is sensitive to spatial distributions of off-axis clumps. $G$ and $M_{20}$ thus should be sensitive to the presence of merger features and indeed LPM04 found that local ULIRGs are well separated from normal galaxies in $G-M_{20}$ space, in the sense that ULIRGs have higher $G$ and $M_{20}$ values than normal galaxies. In this paper we present the results of color-magnitude analysis and $G$ and $M_{20}$ morphological analysis of 54 ULIRGs from the $IRAS$ 1Jy ULIRG sample (Kim 1995; Kim \\& Sanders, 1998) that were observed by Data Release 5 of Sloan Digital Sky Survey. The goals of this paper are to place ULIRGs in the context of normal galaxies by comparing their color-magnitude relations and morphologies with those of optically-selected SDSS galaxies, and to construct a low-redshift comparison sample for studying high-redshift dusty starbursts, such as Spitzer-identified ULIRGs and submillimeter galaxies (SMGs). The $IRAS$ 1Jy ULIRG sample is a complete flux-limited sample from the IRAS Faint Source Catalog (Moshir et el. 1990) with $S_{60\\mu m}>1 \\rm Jy$ and $\\delta>-40^{\\circ}, |b|>30^{\\circ}$ (Kim \\& Sanders 1998). The sample contains 118 sources with a redshift range between 0.018 and 0.268 and a median redshift of 0.145. The infrared luminosities range between $10^{12.00}$ and $10^{12.90}\\rm L_{\\odot}$, with a median of $10^{12.19}\\rm L_{\\odot}$. They are the most IR luminous galaxies in the low redshift universe, and their uniform selection and completeness make them ideal for a statistical analysis. Almost all ULIRGs in the sample show visual signs of interactions or mergers, but most are at advanced stages, and the median R-band luminosity is $\\sim2L_*$ (Kim, Veilleux, \\& Sanders 2002, hereafter KVS02; Veilleux, Kim, \\& Sanders 2002, hereafter VKS02). Spectroscopic study showed that the Seyfert fraction increases with infrared luminosity, and $\\sim50\\%$ of the galaxies with $L_{IR}>10^{12.3}L_{\\odot}$ present Seyfert characteristics. These AGN-powered ULIRGs have bolometric luminosities and near-IR spectra reminiscent of the local quasars. Another $30\\%$ of the whole sample are classified as H{\\sc ii}-region and $\\sim40\\%$ are classified as LINERs, in both of which the energy sources are thought to be stellar origin from recent starbursts ($\\leq$ few times $10^7$ yr) or shocks (Veilleux, Kim \\& Sanders 1999, hereafter VKS99; Kim, Veilleux, and Sanders 1998) Near-IR study using 2MASS showed that Seyfert galaxies among the 1Jy sample have much redder colors and steeper spectral indices in the near-infrared than the rest, which indicates powerful AGNs as the energy sources (Chen \\& Zhang 2006). Recent $^{12}$CO ($J=1\\rightarrow 0$) observations of 17 ULIRGs in the 1Jy Sample, along with 12 other local ULIRGs showed that their large IR luminosity and gas mass to FIR luminosity ratios are consistent with a model where most of their luminosity is powered by a brief surge in star formation rate associated with the rapid gas inflow during the merger phase, although there is also evidence for powerful AGNs with extreme $L_{FIR}/L_{CO}^{\\prime}$ among a subset and they are possibly transitioning to a quasar phase (Chung et al. 2009). In general, previous studies suggest an evolutionary sequence of ULIRGs in which ULIRGs start with cool infrared colors ($f_{25\\mu m}/f_{60\\mu m}<0.2$), go through a warm ULIRG state ($f_{25\\mu m}/f_{60\\mu m}>0.2$), and finally evolve to quasars (Sanders 1988; VKS02). In \\S\\ref{sample_selection} we present the sample selection. The image analysis is presented in \\S\\ref{image_analysis}. We present the color magnitude relation of the ULIRGs in \\S\\ref{cmr}, their morphological analysis in \\S\\ref{morphology}, and the summary in \\S\\ref{summary}. We use a $\\Lambda$CDM cosmology with $H_0=70\\rm kms^{-1}Mpc^{-1}$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$ throughout the paper, except for the absolute magnitudes, for which we use $h=1$, where $h=H_0/100\\rm kms^{-1}Mpc^{-1}$, for a direct comparison with existing studies. ", "conclusions": "Our visual inspection shows that almost all the ULIRGs are disturbed systems (Fig. \\ref{fig2}), and this is consistent with the visual classifications by VKS02, who found that all but one of the 118 $IRAS$ 1Jy sample show visual merger features. We perform the $G-M_{20}$ analysis and find that more disturbed galaxies do have higher $G$ and $M_{20}$ coefficients. This is illustrated in Fig. \\ref{fig1} where more disturbed galaxies tend to occupy the upper left corner of the mosaic with higher values of both $G$ and $M_{20}$. However our studies also show that ULIRGs are a heterogeneous group in $G-M_{20}$ space: only slightly less than half of the sources lie above the solid line in each panel of Fig. \\ref{gm20}, the region where most local mergers and ULIRGs were found by LPM04. Interestingly, there is only one source located in the early type region, and all the remaining sources fall in the region where irregular and late-type galaxies are located. This heterogeneous distribution seems to be inconsistent with the result of LPM04 that 80\\% of local ULIRGs are located in the typical merger regions in the parameter space. However, the spread in the parameter space we found is supported by numerical simulations. Lotz et al. (2008b) analyzed the morphological parameters in SDSS $g-$band images of mergers of equal mass gas-rich spirals by using the N-body/hydrodynamic simulation code GADGET (Springel, Yoshida, \\& White 2001) and Monte Carlo radiative transfer code SUNRISE (Jonsson 2006; Jonsson et al. 2006), and find that the mergers are most disturbed in $G-M_{20}$ at the first pass, % but they can have normal galaxies morphologies at other merger stages. They also noted that two-thirds of the ULIRGs in LPM04 exhibit double or multiple nuclei, % and therefore are more effectively selected by $G-M_{20}$. The timescale is only 0.2 $-$ 0.6 Gyr for a merger to appear in the empirical merger region defined by LPM04, and the range in this time scale depends on dust, orbit parameters, and observing orientations. Thus at face value, our results imply that about half of our ULIRGs have been captured within this time window. The fact that the other half of our ULIRGs do not appear in the empirical merger region does not mean that they are not mergers. Instead, they might be in other merger stages when their morphological parameters are consistent with those of normal galaxies. The heterogeneous picture of our ULIRGs thus may represent different evolutionary epochs of the ULIRGs. In order to gain some insights into the $G-M_{20}$ morphology we compare our $g$-band $G-M_{20}$ classification with the interaction classification for the same sources made by Veilleux et al. (2002). Veilleux et al. classified each object into one of the six sequential merging stages according to its morphology. The comparison of two different classifications is illustrated in Fig. \\ref{morph.comp}. We find that all triple mergers by the classification of Veilleux et al. are also classified as mergers by $G-M_{20}$. Sources in their type IIIa group, i.e., wide binary pre-merger systems with apparent separations greater than 10 kpc, are also very likely to be recognized as $G-M_{20}$ mergers: $\\sim 80\\%$ of such sources are recognized as mergers using $G-M_{20}$. Sources in their other groups are less likely to be recognized by $G-M_{20}$ as mergers, and the fraction of $G-M_{20}$ mergers appears to decrease toward later merging stages. We find that $G-M_{20}$ mergers make up $\\sim 31\\%$ of the close binary pre mergers (type IIIb), $\\sim50\\%$ of the compact mergers (type IVa), $\\sim 46\\%$ of the diffuse mergers (type IVb), and only $\\sim 11\\%$ of the old mergers (type V) are $G-M_{20}$ mergers, respectively. Sources classified as close binary pre-merger systems (type IIIb) and old mergers (type V) have very large fractions to be recognized as spiral galaxies by $G-M_{20}$ ($\\sim 54\\%$ and $\\sim 67\\%$, respectively). These comparison results are qualitatively consistent with the simulation results by Lotz et al. (2008) in which the most disturbed $G-M_{20}$ morphology happens during the first pass and maximum separation. However, since our measured morphology is limited by relatively small number statistics and relatively large uncertainties, we do not conclude strongly that the observed morphology can be well explained by the simulations. We also find that uncertainties in measuring the morphologies contribute significantly to the distribution of the ULIRGs in the parameter space. LPM04 reported that the measurements of $G$ and $M_{20}$ are robust to within 10\\% when the source has signal-to-noise ratio per pixel $\\langle S/N\\rangle $ greater than 2.5, the Petrosian radius larger than 2.5 times the PSF FWHM, and the physical resolution (parsec per pixel) higher than 500 $\\rm pc\\ pixel^{-1}$. However when the physical resolution is lower the uncertainties of morphologies largely increase (Fig. 6 in their paper) since small structures are washed out. At $z\\sim 0.1-0.2$ the spatial resolution of the SDSS ULIRGs is roughly 0.7-1.3 $\\rm kpc\\ pixel^{-1}$ and the uncertainties need to be calibrated carefully. We perform simple simulations to estimate the uncertainties of $G$ and $M_{20}$ for the SDSS ULIRGs at different noise levels. We use four model elliptical galaxies, four spiral galaxies selected from SDSS images, and three merging galaxies selected from our ULIRG sample. The four elliptical galaxies are modeled using de Vacoulers profiles and are added to real SDSS images to mimic the observational conditions. The three spiral galaxies are selected with bright $g$-band magnitudes of 13.9, 15.2, and 16.4, respectively, much brighter than the median $g$-band magnitudes of the SDSS sample (18.0 mag). The three ULIRGs (FSC08572+3915, FSC14060+2919, FSC14121-0126) are selected with bright $g$-band magnitudes and unmistakable merger morphology. The $g$-band magnitudes are 16.4, 16.7 and 17.8, respectively, all brighter than the median g-band magnitudes of the ULIRGs (17.8 mag). For each image we generate random noise at a series of different levels, and at each level we generate 20 noise maps independently and add them to the original image. Morphologies, Petrosian radius, and $\\langle S/N\\rangle$ are measured for each noise-added source and compared to the original measurements. Although our simulation is limited by the sample size, $\\sim90\\%$ of the noise-added sources selected by the signal-to-noise and size criteria ($\\langle S/N\\rangle>2.5$ and $R_p>2.6\\arcsec$) have $g-$band magnitudes between 16 and 19, and $g-$band Petrosian radii between 3 and 12 arcsec ($\\sim$ 6 - 22 kpc at $z\\sim0.1$), ranges similar to those of the ULIRGs. We find that the measured Petrosian radius and the $\\langle S/N\\rangle$ do not always decrease monotonically with the amount of noise added, and they are broadly distributed at low signal-to-noise levels. We also find that, rather than remaining fixed for a given source, the measured $G$ decreases and $M_{20}$ increases systematically with increasing noise. Therefore when seen against higher background noise, a galaxy tends to move to the lower left corner on the $G-M_{20}$ plot. This effect will decrease the fraction of the mergers observed in the merger region. The Gini coefficient $G$ decreases by as much as 0.2 with increasing background noise until the galaxy is no longer detected (either $r_p < 2 \\times FWHM$, or $\\langle S/N\\rangle < 2.5$), and $M_{20}$ increases by as much as 0.5. The standard deviations of $G$ at every noise level are all less than 0.05, and the spiral galaxies have the smallest dispersions. The standard deviations of $M_{20}$ are all less than 0.3-0.4, and the elliptical galaxies have the smallest dispersions. There is no systematic difference for the uncertainties between different bands. These uncertainties ($\\Delta G=0.05$ and $\\Delta M=0.5$) are shown in Fig. \\ref{gm20} as error bars, and the systematic trend with increasing noise in $G$ and $M_{20}$ is shown as an arrow. The aperture within which the morphologies are measured is also an important source of uncertainty. As noted by Lisker (2008), $G$ value measured within larger apertures have systematically larger values, and the measured uncertainties are minimized when the Petrosian radius is used as the aperture size. The reason behind is that more low surface brightness pixels are included within a larger aperture size, which steepens the intensity distribution of the pixels and consequently increases the value of $G$. Some of the low surface brightness features of the source, often at larger distances from the center, are mistakenly excluded when the aperture size is chosen too small, which flattens the pixel intensity distribution and consequently decreases the value of $G$. When assigning pixels to the sources we adopt the surface brightness at the Petrosian radius as the threshold, and assign pixels brighter than the threshold to the source. Therefore, according to their study our measured morphologies should have minimized uncertainties, because our apertures should be very close to the Petrosian radius and not significantly larger or smaller. \\subsection{Relations among Morphology, Optical Color, and FIR Luminosity} The morphology, color-magnitude relation, and the infrared luminosity of ULIRGs present a rather complicated picture of these systems. The huge infrared luminosity in our ULIRGs suggests that their dominant powering sources should be obscured by dust. However, the majority of them are optically bright and blue. This further implies that dust is not uniformly distributed and that significant amount of the unobscured stellar light is seen directly. Indeed, as shown in Fig. 1, there are obvious color gradients in the ULIRGs, with their tidal features at large distances appearing blue and their central regions appearing red. Lotz et al. (2008b) suggested that merging systems are most disturbed in $G-M_{20}$ space and are located in the merger region in the $G-M_{20}$ plot during the first pass when the tidal features are prominent. % After the first pass they will move to the late-type and irregular region (see Fig. 5 in their paper). Thus we would expect the integrated colors of those disturbed systems selected by $G-M_{20}$ to be relatively bluer than the ones with normal $G-M_{20}$ relations, due to the blue colors of the tidal features. Recent hydrodynamic simulations of disc-disc mergers also suggested that the infrared luminosity of ULIRGs will peak during the final merger when enough metals and dust have been accumulated (e.g. Jonsson et al. 2006), so one would also expect sources with higher infrared luminosity are preferentially found in the late-type and irregular region. We test these hypotheses by plotting the $g-$band $Gini$ vs $M_{20}$ for the ULIRGs in Fig. \\ref{other.relation}, with different symbol sizes representing their FIR luminosities and different symbol colors representing their $^{0.1}g-^{0.1}r$ colors. There are no strong relations between the optical colors and $G$ or $M_{20}$, except that optically redder sources are slightly more preferentially located in the non-merger regions with lower $G$ values. This is further illustrated in the small panel in the same plot, where we plot the color distributions of sources with $G\\ge0.55$, the median $G$ of the whole sample, and of sources with $G<0.55$. The sources with lower $G$ values tend to distribute toward redder optical colors, although the reddest source has a higher-than-median $G$ value (0.62). All AGN ULIRGs have $G$ values greater than 0.5. There are no obvious correlations between the FIR luminosity and $G$ or $M_{20}$. These results are broadly consistent with the scenario that the blue optical colors originate from more disturbed systems, and these systems are not during the phase when the FIR emission peaks. However, as mentioned earlier, the uncertainties in measuring $G$ and $M_{20}$ could be substantial, and the measured morphologies are affected by orientation, orbital parameters, viewing angle and dust content (Lotz et al. 2008b). We thus do not conclude any strong relations among the morphologies, the FIR luminosity, and the optical color of these system." }, "1002/1002.2013_arXiv.txt": { "abstract": "{} { We compare six popularly used evolutionary population synthesis (EPS) models through fitting the full optical spectra of six representative types of galaxies (star-forming and composite galaxies, Seyfert 2s, LINERs, E+A and early-type galaxies), which are taken from the Sloan Digital Sky Survey (SDSS); and we also explore the dependence of stellar population synthesis results on the main ingredients of the EPS models; meanwhile we study whether there is an age sequence among these types of galaxies. } {Throughout our paper, we use the simple stellar populations (SSPs) from each EPS model and the software STARLIGHT to do our fits. Firstly, to explore the dependence of stellar population synthesis on EPS models, we fix the age, metallicity, and initial mass function (IMF) to construct a standard SSP group. We then use the standard SSP group from each EPS model (BC03, CB07, Ma05, GALEV, GRASIL, and Vazdekis/Miles) to fit the spectra of star-forming and E+A galaxies. Secondly, we fix the IMF and change the selection of age and metallicity respectively to construct eight more SSP groups. Then we use these SSP groups to fit the spectra of star-forming and E+A galaxies to check the effect of age and metallicity on stellar populations. Finally, we also study the effect of stellar evolution track and stellar spectral library on our results. At the same time, the possible age sequence among these types of galaxies are suggested. } {Using different EPS models the resulted numerical values of contributed light fractions change obviously, even though the dominant populations are consistent. The stellar population synthesis does depend on the selection of age and metallicity, while it does not depend on the stellar evolution track much. The importance of young populations decreases from star-forming, composite, Seyfert 2, LINER to early-type galaxies, and E+A galaxies lie between composite galaxies and Seyfert 2s in most cases. } {Different EPS models do derive different stellar populations, so that it is not reasonable to directly compare stellar populations estimated from different EPS models. To get reliable results, we should use the same EPS model for the compared samples.} ", "introduction": "\\label{introduction} Stellar populations are fundamental characters in revealing the formation and evolution of galaxies. The formation of spiral galaxies is still hotly debated, and two main channels have been proposed, either initial collapse of gas at very high redshift in the frame of the tidal torque theory (Eggen, Lynden-Bell \\& Sandage, 1962; White, 1984), or gaseous-rich mergers at intermediate to high redshifts (Hammer et al., 2005, 2007, 2009). These two channels for galaxy formation may provide distinct signatures from the analysis of the stellar populations, and it is relevant to test whether or not stellar population models can be used to test them (see for example Heavens et al., 2004; Panter et al, 2007). While stars can not be resolved for a majority of galaxies, therefore many works have been generated on analyzing stellar populations through the integrated lights of the galaxies. Because the integrated lights hold information about age and metallicity distributions of their stellar populations and star formation histories. This is the so-called stellar population synthesis on galaxies. Two main types of approaches have been developed: the empirical population synthesis (Faber 1972; Bica 1988; Boisson et al. 2000; Cid Fernandes et al. 2001) and the EPS ( Tinsley 1978; Bruzual 1983; Worthey 1994; Leitherer \\& Heckman 1995; Maraston 1998; Vazdekis \\& Arimoto 1999; Bruzual \\& Charlot 2003; Maraston 2005; Cid Fernandes et al. 2005). In the empirical population synthesis approach, also known as $'stellar\\,population\\,synthesis\\, with\\, a\\, data\\, base'$, the observed spectrum of a galaxy is reproduced by a combination of spectra of individual stars or star clusters with different ages and metallicities from a library. The results following this approach do not consider the stellar evolution, and do not allow one to predict the past and future spectral appearance of galaxies. The EPS approach uses the knowledge of stellar evolution to model the spectrophotometric properties of stellar populations, and has enjoyed more widespread use recently. In this approach, the main adjustable parameters are the stellar evolution tracks, the stellar spectral library, the IMF, the star formation history (SFH), and the grids of ages and metallicities. EPS is a real physical model, but it is restricted by the lacking of comprehensive stellar spectral library, accurate IMF and SFH, and poor understanding of some advanced phases of stellar evolution, such as the blue stragglers (BSs), the horizontal branch (HB) stars, and the thermally pulsating asymptotic giant branch (TP-AGB) stars. Up to now, several EPS models have been proposed and widely used in the stellar population studies on galaxies by analyzing their colors, spectra and multi-wavelength spectral energy distributions (SEDs), such as BC93 (Bruzual \\& Charlot 1993), P\\'{e}gase (Fioc \\& Rocca-Volmerange 1997), GRASIL (Silva et al. 1998), GALAXEV (Bruzual \\& Charlot 2003, BC03), CB07 (Charlot \\& Bruzual 2009), SPEED (Jimenez et al. 2004), BaSTI (Pietrinferni et al. 2004, Cordier et al. 2007), Ma05 (Maraston 2005), Starburst 99 (V\\'{a}zquez \\& Leitherer 2005), Vazdekis/Miles (S\\'{a}nchez-Bl\\'{a}zquez et al. 2006; Vazdekis et al., in preparation), GALEV (Anders \\& Alvensleben 2003; Kotulla et al. 2009), SPoT (Raimondo et al. 2005). Some researches have used these EPS models to analyze the stellar populations in galaxies and star clusters, and even compared them. The results are interesting, however, most of them focus on analyzing the colors and multi-wavelength SEDs of the systems. Maraston et al. (2006) used two sets of EPS models to estimate the star formation histories, ages, and masses of seven galaxies in the Hubble Ultra Deep Field by analyzing their observed spitzer mid-IR (the rest-frame-UV) photometry data. One of the EPS models is Ma05, which includes the contribution of TP-AGB stars, and another one is represented by BC03 (similar models are P\\'{e}gase, Starburst 99 etc.). They concluded when they assumed a zero reddening, Ma05 gave better fits than BC03 for these distant passively evolving galaxies at $1.4 3$, implying spatially anisotropic fluctuations, $k_{\\perp} > k_{\\para}$. The spectral index of the perpendicular component is $-2.6$ at large angles and $-3$ at small angles, in broad agreement with critically balanced whistler and kinetic \\Alfven\\ wave predictions. For the parallel component, however, it is shallower than $-1.9$, which is considerably less steep than predicted for a kinetic \\Alfven\\ wave cascade. ", "introduction": " ", "conclusions": "" }, "1002/1002.2155_arXiv.txt": { "abstract": "{The Kolmogorov stochasticity parameter is shown to act as a tool to detect point sources in the cosmic microwave background (CMB) radiation temperature maps. Kolmogorov CMB map constructed for the WMAP's 7-year datasets reveals tiny structures which in part coincide with point radio and Fermi/LAT gamma-ray sources. In the first application of this method, we identified several sources not present in the then available 0FGL Fermi catalog. Subsequently they were confirmed in the more recent and more complete 1FGL catalog, thus strengthening the evidence for the power of this methodology.} \\begin{document} ", "introduction": "A number of point sources are identified in the temperature maps of cosmic microwave background (CMB), including in the Wilkinson Microwave Anisotropy Probe's (WMAP) data \\cite{Ja}; the point source catalog is available in http://lambda.gsfc.nasa.gov/product/map/. The efficiency of the search of point sources is particularly important for obtaining a pure cosmological signal of high precision at large multipoles, for the cross correlation studies with galaxy surveys, on baryonic oscillations, e.g. \\cite{Saw1}\\cite{Saw2}. We now involve for this aim a new descriptor, the Kolmogorov stochasticity parameter which describes the degree of randomness of a given sequence \\cite{Kolm}\\cite{Arnold}. It has been already been applied to the CMB temperature datasets \\cite{KSP}. The resulting Kolmogorov CMB maps showed the potential of this approach in the separation of the Galactic disk from the cosmological signal, in the identification of anomalies such as the non-Guassian Cold Spot \\cite{Cruz}\\cite{anomaly4}. Certain spots and regions noticed in the K-map can be compared with the studies with other descriptors \\cite{Ross}. Here we constructed and analyzed Kolmogorov's map for WMAP's 7-year temperature data\\footnote{This study initially was based on the WMAP's 5-year data and 7-year data were involved later upon their release: the results are coinciding absolutely.} and obtained a list of 12 regions of a degree scale with anomalously high value of Kolmogorov's function (see below). Among those regions, 6 coincide with sources given in the catalog of the WMAP's point sources, where they are identified as radio galaxies, while all 12 regions coincide with gamma-ray sources discovered by Fermi satellite's Large Area Telescope (LAT) \\cite{Fermi1}; 4 among these regions in both catalogs are overlapping mutually. In the first version of this study (arXiv:1002.2155v), when only 0FGL catalog of Fermi/LAT was available \\cite{Fermi0}, there were 3 regions which have practically identical values both of the Kolmogorov's function and of the CMB temperature distribution but were absent in either of mentioned catalogs, i.e. of the radio and Fermi gamma sources. One might thus expect that they also are still unidentified sample of point sources. Indeed, the release of Fermi/LAT 1FGL catalog confirmed this prediction, i.e. all 12 CMB structures are identified in that catalog, thus showing the predictive power of the approach. The detected sources include blazars/flat spectrum radio quasars (FSRQs) and active galaxy \\cite{Fermi1}, which are distinguished by their multi-wavelength emission, from radio to gamma rays, and due to their superluminous nature are efficient in probing the distant Universe. ", "conclusions": "We used the Kolmogorov's parameter to extract point sources in the CMB maps. The idea is that the radiation of different origin can have different statistic and this can be revealed by the Kolmogorov's distribution. Numerical experiments for systems with built a priori different statistic showed that K-parameter is indeed sensitive to the randomness properties of systems \\cite{mod}. Namely, the regions which might be of lower significance by conventional methods and hence are absent in WMAP's catalog, can be easily outlined by the present approach. On the other hand, the increase in the angular resolution of the CMB temperature maps will enable to apply this method for smaller scales, i.e. to obtain the averaged Kolmogorov function for large enough number of pixels, and hence to increase its source detecting power. The analysis performed for WMAP's 7-year w-band maps aimed the extraction of around 1 degree scale regions (namely, when averaged within $1.0^{\\circ}$ and $1.5^{\\circ}$ circles at $0.2^{\\circ}$ and $0.5^{\\circ}$ step scans, respectively), which were outlined in the K-map due to their high value of the Kolmogorov's function $\\Phi$. When compared with the catalog of the point sources of WMAP's corresponding map, 6 among those $\\Phi$-regions were identified with radio galaxies, while all 12 regions were identified with Fermi/LAT gamma-ray 1FGL sources \\cite{Fermi1}. The 3 detected anomal regions were absent in the previous 0FGL Fermi/LAT catalog \\cite{Fermi0}. However as concluded then, since the CMB regions have practically identical distributions of the function $\\Phi$ as well as of the CMB temperature, the latter regions can also be point sources of close nature. Present study confirmed that prediction and the interest for dedicated studies of CMB map's those regions in radio (see e.g. \\cite{J}) and other bands. At the available angular resolution of WMAP the method enables to identify at high confidence level only a sample of 12 regions, which is certainly smaller than the number of sources identified by WMAP's mask. Again, the reason, as described above, is in the available resolution of the WMAP's maps. However, even at present resolution the method identifies sources not noticed by WMAP's procedure, which appear to be Fermi/LAT gamma sources \\cite{Fermi1}. Most of those sources appear to be blazar/quasars and active galaxy, i.e. multi-wavelength and superluminous objects, and we showed another possibility of their identification based on totally different idea, the statistic of the signal. The absence of some of the detected sources in the initial 0FGL catalog and their appearance in the next, 1FGL catalog, shows the predictive power of the approach. The Kolmogorov parameter thus enables to detect point sources in the CMB maps not noticed by other criteria. The applicability of the method will increase at higher resolution maps expected from the PLANCK mission. \\begin{table*} { \\renewcommand{\\baselinestretch}{1.5} \\renewcommand{\\tabcolsep}{1.3mm} \\small{ {\\bf Table 1.} Coordinates of the 12 regions in WMAP's CMB map outlined by Kolmogorov's function $\\Phi$, along with the coordinates of the identified radio sources, with the accuracy of identification in arcminutes dA and source designations, and the same for Fermi/LAT gamma-ray sources. \\\\[5pt] } \\small{ \\begin{center} \\begin{tabular}{crrrrcccrccl} \\hline \\multicolumn{3}{c}{$\\Phi$}& \\multicolumn{4}{c}{Radio} &\\multicolumn{4}{c}{Fermi/LAT}\\\\ No&\\multicolumn{1}{c}{{\\it l}}&\\multicolumn{1}{c}{{\\it b}} &\\multicolumn{1}{c}{{\\it l}}&\\multicolumn{1}{c}{{\\it b}}& D{\\it A}& &\\multicolumn{1}{c}{{\\it l}}&\\multicolumn{1}{c}{{\\it b}}& dA& 1FGL&nature \\cite{Fermi1} \\\\[2pt] \\hline \\hline 01&283.55& 74.55&283.81& 74.49&5.50&GB6 J1230+1223&283.83& 74.50& 5.36&J1230.8+1223 &active galaxy \\\\ 02&289.90& 64.40&289.95& 64.36&2.73&GB6 J1229+0202&289.95& 64.36& 2.70&J1229.1+0203 &FSRQ \\\\ 03&305.10& 57.10&305.11& 57.06&2.42&PMN J1256-0547&305.12& 57.06& 2.44&J1256.2-0547 &FSRQ \\\\ 04&309.50& 19.40& & & & &309.55& 19.42& 3.34&J1325.6-4300 &active, radio galaxy \\\\ 05&353.00& 16.80& & & & &353.47& 16.58&30.33&J1628.6-2419c&\\\\ 06&195.35&-33.15&195.29&-33.14&3.07&PMN J0423-0120&195.25&-33.14& 5.30&J0423.2-0118 &FSRQ\\\\ 07&206.50&-16.35& & & & &206.72&-16.38&12.82&J0541.9-0204c&\\\\ 08&208.90&-19.50& & & & &209.07&-19.56&10.43&J0534.7-0531c&\\\\ 09&213.70&-12.60& & & & &213.86&-12.58& 9.52&J0608.1-0630c&\\\\ 10&279.60&-31.60& & & & &279.62&-31.63& 2.38&J0538.9-6914 &normal galaxy\\\\ 11& 9.40&-19.60& 9.35&-19.61&2.89&PMN J1924-2914& 9.31&-19.72& 9.11&J1925.2-2919 &FSRQ\\\\ 12& 86.00&-38.20& 86.12&-38.19&5.69&GB6 J2253+1608& 86.12&-38.19& 5.55&J2253.9+1608 &FSRQ\\\\ \\hline \\end{tabular} \\end{center} }} \\end{table*}" }, "1002/1002.0366_arXiv.txt": { "abstract": "The air fluorescence detector of the \\pao is designed to perform calorimetric measurements of extensive air showers created by cosmic rays of above $10^{18}$~eV. To correct these measurements for the effects introduced by atmospheric fluctuations, the Observatory contains a group of monitoring instruments to record atmospheric conditions across the detector site, an area exceeding 3,000 km$^2$. The atmospheric data are used extensively in the reconstruction of air showers, and are particularly important for the correct determination of shower energies and the depths of shower maxima. This paper contains a summary of the molecular and aerosol conditions measured at the \\pao since the start of regular operations in 2004, and includes a discussion of the impact of these measurements on air shower reconstructions. Between $10^{18}$ and $10^{20}$~eV, the systematic uncertainties due to all atmospheric effects increase from $4\\%$ to $8\\%$ in measurements of shower energy, and $4$~\\gcmsq to $8$~\\gcmsq in measurements of the shower maximum. ", "introduction": "\\label{sec:introduction} The \\pao in \\mal, Argentina (69$^{\\circ}$ W, 35$^{\\circ}$ S, 1400~m a.s.l.) is a facility for the study of ultra-high energy cosmic rays. These are primarily protons and nuclei with energies above $10^{18}$~eV. Due to the extremely low flux of high-energy cosmic rays at Earth, the direct detection of such particles is impractical; but when cosmic rays enter the atmosphere, they produce extensive air showers of secondary particles. Using the atmosphere as the detector volume, the air showers can be recorded and used to reconstruct the energies, arrival directions, and nuclear mass composition of primary cosmic ray particles. However, the constantly changing properties of the atmosphere pose unique challenges for cosmic ray measurements. In this paper, we describe the atmospheric monitoring data recorded at the \\pao and their effect on the reconstruction of air showers. The paper is organized as follows: Section~\\ref{sec:atmocal} contains a review of the observation of air showers by their ultraviolet light emission, and includes a description of the \\pao and the issues of light production and transmission that arise when using the atmosphere to make cosmic ray measurements. The specifics of light attenuation by aerosols and molecules are described in Section~\\ref{sec:prod_and_trans}. An overview of local molecular measurements is given in Section~\\ref{sec:molecular_effects}, and in Section~\\ref{sec:measurements} we discuss cloud-free aerosol measurements performed at the Observatory. The impact of these atmospheric measurements on the reconstruction of air showers is explored in Section~\\ref{sec:aerosol_effects}. Cloud measurements with infrared cameras and backscatter lidars are briefly described in Section~\\ref{sec:future}. Conclusions are given in Section~\\ref{sec:conclusion}. \\input atmocal \\input atmotrans \\input measmol \\input measaer \\input hybridres \\input futuredevel ", "conclusions": "\\label{sec:conclusion} A large collection of atmospheric monitors is operated at the \\pao to provide frequent observations of molecular and aerosol conditions across the detector. These data are used to estimate light scattering losses between air showers and the FD telescopes, to correct air shower light production for various weather effects, and to prevent cloud-obscured data from distorting estimates of the shower energies, shower maxima, and the detector aperture. In this paper, we have described the various light production and transmission effects due to molecules and aerosols. These effects have been converted into uncertainties in the hybrid reconstruction. Most of the reported uncertainties are systematic, not only due to the use of local empirical models to describe the atmosphere --- such as the monthly molecular profiles --- but also because of the nature of the atmospheric uncertainties --- such as the systematics-dominated and highly correlated aerosol optical depth profiles. Molecular measurements are vital for the proper determination of light production in air showers, and molecular scattering is the dominant term in the description of atmospheric light propagation. However, the time variations in molecular scattering conditions are small relative to variations in the aerosol component. The inherent variability in aerosol conditions can have a significant impact on the data if aerosol measurements are not incorporated into the reconstruction. Because the highest energy air showers are viewed at low elevation angles and through long distances in the aerosol boundary layer, aerosol effects become increasingly important at high energies. Efforts are currently underway to reduce the systematic uncertainties due to the atmosphere, with particularly close attention paid to the uncertainties in energy and \\xmax. The shoot-the-shower program will improve the time resolution of atmospheric measurements, and increase the identification of atmospheric inhomogeneities that can affect observations of showers with the FD telescopes." }, "1002/1002.1082_arXiv.txt": { "abstract": "Gravitational waves (GW) can be emitted from coalescing neutron star (NS) and black hole-neutron star (BH-NS) binaries, which are thought to be the sources of short hard gamma ray bursts (SHBs). The gamma ray fireballs seem to be beamed into a small solid angle and therefore only a fraction of detectable GW events is expected to be observationally coincident with SHBs. Similarly ultrahigh energy (UHE) neutrino signals associated with gamma ray bursts (GRBs) could fail to be corroborated by prompt $\\gamma$-ray emission if the latter is beamed in a narrower cone than the neutrinos. Alternative ways to corroborate non-electromagnetic signals from coalescing neutron stars are therefore all the more desirable. It is noted here that the extended X-ray tails (XRT) of SHBs are similar to X-ray flashes (XRFs), and that both can be attributed to an off-axis line of sight and thus span a larger solid angle than the hard emission. It is proposed that a higher fraction of detectable GW events may be coincident with XRF/XRT than with hard $\\gamma$-rays, thereby enhancing the possibility to detect it as a GW or neutrino source. Scattered $\\gamma$-rays, which may subtend a much larger solid angle that the primary gamma ray jet, are also candidates for corroborating non-electromagnetic signals. ", "introduction": "Short hard $\\gamma$- ray bursts (SHBs) are now suspected to be caused by the merging of two compact objects, such as neutron stars or black holes, which would release large amounts of energy over short time intervals (e.g. Goodman, 1986). Collapse of a single object has also been proposed to give rise to a similar situation (Berezinsky 1987). Eichler et al (1989) suggested that GRBs could be observed in coincidence with GW signals when two neutron stars merge. The huge isotropic equivalent energy requirements implied by the BATSE observations of GRB isotropy and submaximal $V/V_{max}$, suggested that GRBs might be highly collimated (e.g. Levinson and Eichler, 1993) and this would make them a bad bet to corroborate GW signals from such mergers, as gravitational radiation is unlikely to be strongly collimated. This might be fatal (e.g. Guetta and Stella, 2009) to the original proposal that LIGO signals would be coincident with GRBs. It is now accepted that GRBs are indeed highly collimated. Alternative ways to confirm LIGO signals from coalescing neutron stars [and, according some suggestions (Van van Putten 1999a,b), unstable collapsing disks] are therefore all the more desirable. The horizons of first generation LIGO and Virgo for NS-NS, and (BH-NS) mergers are $\\sim 20$ and $43$ Mpc, respectively, while advanced LIGO/Virgo should detect them out to a distance of $\\sim 300$ and $650$ Mpc (for a review see Cutler \\& Thorne 2002). GW signals from NS-NS mergers are expected at a rate of one in 10-150 years with Virgo and LIGO and one every 1-15 days with Advanced LIGO/Virgo class interferometers (Berezinsky et al. 2002, 2007). The BH-NS and BH-BH merger rates in the Galaxy are highly uncertain. Berezinsky et al. (2007) estimate 1\\% and 0.1\\% of the NS-NS merger rate, respectively, implying that BH-NS and BH-BH mergers contribute marginally to the GW event rate, despite the greater distance out to which they can be detected. Ultahigh energy (UHE) neutrinos may come from nearby supernova even if an associated GRB is shaded from our view or entirely smothered by the envelope of the host star (Eichler and Levinson, 1999). A fluence $F$ of $10^{-4}$erg/cm$^2$ in muon neutrinos at $10^{12} \\le E_{\\nu} \\le 10^{14}$ eV yields roughly a single neutrino detection in a gigaton detector such as ICECUBE, the exact expectation value depending somewhat upon the energy. An UHE neutrino signal from a nearby supernova or supernova/GRB could therefore be detectable at a distance of $D\\sim 10^2 E_{iso,\\nu,50}^{1/2}$ Mpc. Note that $E_{iso}$ for $\\gamma$-rays can be as high as $10^{54}$ erg [$E_{iso,\\gamma,50} = 10^4$]. We face the following interesting possibility: If the UHE neutrinos from GRB are beamed into a wider beam than the $\\gamma$-rays, then even if the neutrino efficiency is high, the value of $E_{iso,\\nu}$ may be too low to be seen from any given burst unless it is close. More importantly, most UHE neutrino events from GRB sources would not coincide with observed GRBs, as the latter would be most likely beamed away from us. For corroborating UHE neutrino signals, as is the case for GW corroboration, we therefore seek electromagnetic signals that have broader angular reach than the primary $\\gamma$-rays, even at the expense of $E_{iso}$. We note that several wide angle manifestations of nearby GRBs have been proposed. Eichler and Levinson (1999) have suggested both high energy neutrino signals and scattered photons [i.e. scattered off material moving at Lorentz factor less than the intrinsic opening angle of the primary emission, see also Eichler and Manis (2007)], each of which could corroborate LIGO events, at large viewing offsets. Levinson et al (2002) have suggested orphan afterglows, though there might be some problem establishing uniqueness via coincidence because of their long timescales. As it happens, evidence for a high degree of collimation is more convincing for the long GRBs, while SHBs are the ones now believed to be associated with mergers. SHBs frequently show a lower $E_{iso}$, a somewhat broader $V/V_{max}$ distribution, and less evidence of a narrow opening angle from jet breaks. This could be understood, for example, if the giant envelope in the case of long bursts provides better collimation than when it is absent. The presence or absence of the envelope may be responsible for other differences between short and long GRBs. For example, it may be that there is intrinsic spread in the timescales of the central engines accretion timescale, and that only long bursts are sustained enough to break through a massive envelope, whereas mergers, perhaps for different reasons, also produce a spread in timescales while allowing all of them to be observed, though this would explain neither spectral differences nor differences in spectral lags and sub-pulse time scales between short and long GRBs. Neither would it by itself explain why the short duration, hard emission of short GRB lie off to one side of the Amati relation while long bursts, X-ray flashes and the X-ray tails of short GRB obey it. Eichler and Manis (2007, 2008) and Eichler, Guetta and Manis (EGM, 2009) noted that the unusually hard spectrum displayed by SHBs, their unusually soft X-ray tail (as compared to the emission of long GRBs), and their short duration were consistent with a smaller Lorentz factor at the time the short, hard emission is last scattered, and a larger viewing angle. The larger viewing angle is, {\\it a priori}, statistically expected if the line of sight is not obscured by an extended stellar envelope which is known to exist in the case of long GRBs, and which would be less likely in the case of NS-NS mergers. Less collimation and larger allowed viewing offset angle make a coincidence with a GW signal more likely. While larger viewing angle and/or less collimation means smaller $E_{iso}$ and therefore less $V_{max}$, that is not a problem for LIGO collaboration, where the sources would in any case be very close. \\footnote{The suggestion of Eichler and Manis (2007, 2008) and EGM (2008) that viewing angle affects the perceived durations both of the SHBs hard emission and X-ray tail is compatible with an additional intrinsic spread in durations of central engine activity for mergers (van Putten 1999b). The long duration of GRB060614 can be attributed to the prolonged activity of a rotating black hole (Van Putten, 2008). The hypothesis of EGM can also accommodates events such as GRB 060614, which was of long duration while resembling SHBs in other respects. Also, an X-ray tail that lasts $10^2$ s in observer time can result from a SHB whose intrinsic duration is only 1 s. In this paper, however, we are concerned only with the angular spread of the X-ray tail, not the intrinsic duration of the central engine activity that causes it, and consider the possibility that the X-ray tails of SHBs may have broader angular spreads and encompass more observers than the short, hard emission.} Admittedly, the typical viewing angle for SHBs, though perhaps larger than for long GRBs, is uncertain and could be small compared to unity. There exists by now some evidence that SHBs are beamed, like long GRBs, into a modest solid angle. Fox et al. (2005) interpreted the steepening of the optical afterglow light curve of GRB 050709 and GRB 050724 in terms of a jet break that translates into a beaming factor $f_b^{-1}\\sim 50$ (with $f_b$ the fraction of the $4\\pi$ solid angle within which the GRB is emitted). Soderberg et al (2006) found a beaming factor of $\\sim 130$ for GRB 051221A. Therefore with the present data the beaming angles of SHB seem to lie in the range of $\\sim 0.1-0.2$ radian. The discovery that X-ray flashes (XRFs) are a class of long GRBs was made by the Wide Field Camera (WFC) on BeppoSax (Heise et al. 2001). The XRFs are GRBs characterized by no or faint signal in the $\\gamma$-ray energy range, are isotropically distributed in the sky and have an average duration of $\\sim 100$ sec like long GRBs. There is strong evidence that classical GRBs and XRFs are closely related phenomena, and understanding what makes them differ could yield important insights into their origin. The redshift distribution of XRFs is very similar to normal GRBs and therefore the high redshift hypothesis, which might otherwise justify the softness of the burst, cannot account for all XRFs. D'Alessio et al. (2006) have concluded that the off-axis hypothesis seems to be the best hypothesis for now. Many SHBs show a bright X-ray tail (XRT) that follows the short prompt $\\gamma$-ray emission and lasts for $\\sim 100$ s (e.g. Nakar 2007). This X-ray component is evident in SHBs 050709 and 050724, where the X-ray energy is comparable to or even larger than the energy in the prompt $\\gamma$-rays. It seems that XRTs, though not detected in all SHBs, are rather common among them. Extrapolation of the late afterglow back to early times suggests that these tails cannot be interpreted as the onset of the afterglow emission (Nakar 2007). These X-ray tails have spectra and durations that are similar to those of the know XRFs, and maybe both can be attributed to an off-axis line of sight. In this case, they could encompass more observers than the hard emission of the SHB, and could thus be more opportune for corroborating non-electromagnetic manifestations of mergers and/or core collapses. EGM (2009) made rough estimates of order 0.1 to 0.2 radian offset from the periphery of the primary fireball, but with large uncertainties. In this paper we consider the possibility that a wider opening angle of X-ray tails, relative to the hard SHB emission, enhances their likelihood of corroborating non-electromagnetic signals from merger and collapse events. In section 2 we report all the properties of the X-ray tails. In section 3 we present a method to determine the XRFs rate. In section 4 we compare this rate with the XRT rate and discuss our results. ", "conclusions": "The rate of XRFs is an upper limit to the rate of XRTs. For the lower limit we can take the one of SHBs derived by Guetta and Stella (2009). In this paper they find evidence in favor of a bimodal origin of SHB progenitors where a fraction of SHBs comes from the merging of primordial neutron star-neutron star (black hole) and a fraction comes from the merging of dynamically formed binaries in galaxy clusters. In particular they find that the redshift distribution of SHBs is best fitted when the incidences of primordial and dynamical mergers among the SHB population are 40\\% and 60\\% respectively. In this case the rate of SHB is $R_0\\sim 2.9$~Gpc$^{-3}$yr$^{-1}$. For a fiducial value of $f_b^{-1}\\sim 100$, we derive a beaming-corrected rate of $\\rho_0 =f_b^{-1} R_0\\sim 290(f_b^{-1}/100)$. Therefore the rate of XRTs is $2.91$ to 5 (e.g., Brown \\& Vanden Bout 1992; Omont et al. 1996a; Carilli et al. 2002; Cox et al. 2002; Solomon \\& Vanden Bout 2005, hereafter SV05; Riechers et al. 2006; Maiolino et al. 2007; Coppin et al. 2008a). These CO detected quasars show a FIR-to-CO luminosity correlation similar to local starburst spiral galaxies, Ultra Luminous Infra Red Galaxies (ULIRGs), and high-z submillimeter galaxies (SMGs, SV05; Greve et al. 2005; Riechers et al. 2006). The derived molecular gas masses are on the order of $\\rm \\sim10^{10}\\,M_{\\odot}$, which are also comparable to the typical values found in SMGs (Greve et al. 2005; Carilli \\& Wang 2006; Coppin et al. 2008a) and the z$\\sim$1.5 massive star forming disk galaxies (Daddi et al. 2008). The molecular gas can provide the requisite fuel for massive star formation which is suggested by the FIR emission redshifted to submillimeter and millimeter wavelengths (Benford et al. 1999; Omont et al. 1996b, 2003; Priddey et al. 2003; Robson et al. 2004; Beelen et al. 2006; Wang et al. 2008a). The CO line emission detected in FIR luminous quasars provides estimates of the dynamical properties of the spheroidal bulges (i.e. bulge mass and velocity dispersion). Shields et al. (2006) investigated the relation between black hole mass ($\\rm M_{BH}$) and bulge velocity dispersion ($\\rm \\sigma$) in high-z quasar systems using the observed CO line widths. They found that the massive quasars ($\\rm M_{BH}>10^{9}$ to $\\rm 10^{10}\\,M_{\\odot}$) at z$>$3 appear to have much narrower $\\rm \\sigma$ values compared to what is expected from the local $\\rm M_{BH}$--$\\rm\\sigma$ relationship (e.g., Tremaine et al. 2002). Moreover, Coppin et al. (2008a) found that the average black hole--bulge mass ratios for CO and FIR luminous quasar samples at z$\\sim$2 are likely to be an order of magnitude higher than the local value ($\\rm M_{BH}/M_{bulge}\\sim0.0014$, Marconi \\& Hunt 2003). Less evolved stellar bulges were also indicated by high resolution imaging of the CO emission in two z$>$4 quasars (Riechers et al. 2008a, 2008b). These results argue for the scenario that the formation of the SMBHs occurs prior to that of the stellar bulge, which is also suggested by the near-IR imaging of high-z quasar host galaxies (Peng et al. 2006a, 2006b). CO line emission was previously searched in four z$\\sim$6 quasars, and detected in two of these (Walter et al. 2003; Bertoldi et al. 2003b; Carilli et al. 2007; Maiolino et al. 2007). The two CO-detected quasars, J1148+5251 and J0927+2001, are the brightest millimeter sources among a sample of thirty-three z$\\sim$6 quasars that have published millimeter dust continuum observations (Petric et al. 2003; Bertoldi et al. 2003a; Wang et al. 2007, 2008a). The detected CO transitions indicate highly excited molecular gas in the quasar host galaxies (line flux density spectral energy distributions peaking at \\textit{J}=6 or higher, Carilli et al. 2007; Riechers et al. 2009) and the implied molecular gas masses are all $\\rm \\sim2\\times10^{10}\\,M_{\\odot}$. High resolution VLA imaging has resolved the CO emission in J1148+5251 to kpc scale, suggesting a dynamical mass of $\\rm M_{dyn}sin^2{\\it i}\\sim4.5\\times10^{10}\\,M_{\\odot}$ within a radius of 2.5 kpc, where {\\it i} is the inclination angle (Walter et al. 2004; Riechers et al. 2009) of the molecular disk. This provides the first direct constraint on the mass of the quasar host galaxy at the highest redshift, which suggests a high black hole--bulge mass ratio similar to that found with the z$\\sim$2 and z$>$4 quasars. In this work, we extend the CO observations to all z$\\sim$6 quasars with known 250 GHz continuum flux densities of $\\rm S_{250GHz}\\geq1.8\\,mJy$ (Wang et al. 2007, 2008b). We aim to study their general molecular gas properties and investigate possible constraints on the SMBH-host evolution with these earliest quasars. We describe the sample selection and observations in Section 2, and present the results in Section 3. The CO emission and gas properties of these z$\\sim$6 quasars are analyzed in Section 4. Based on our results, we present a brief discussion on the constraints of the black hole-bulge evolution in Section 5, and summarize the main conclusions in Section 6. A $\\rm \\Lambda$-CDM cosmology with $\\rm H_{0}=71km\\ s^{-1}\\ Mpc^{-1}$, $\\rm \\Omega_{M}=0.27$ and $\\rm \\Omega_{\\Lambda}=0.73$ is adopted throughout this paper (Spergel et al. 2007). ", "conclusions": "We present new observations of molecular CO line emission and 350 $\\mu$m dust continuum emission in quasar host galaxies at z$\\sim$6. Our most important finding is that high-order CO transitions are detected in all six of the z$\\sim$6 quasars observed with the 3 mm receiver on the PdBI. The new CO detections all have observed 250 GHz dust continuum flux densities of $\\rm S_{250GHz}\\geq1.8\\,mJy$. These results, together with previous CO-detections in another two objects, reveal an extremely high CO detection rate in the FIR luminous quasars at z$\\sim$6. With the final sample of eight CO-detected z$\\sim$6 quasars, we study the molecular gas properties in the earliest quasar host galaxies, and the main results are summarized as follows: The CO emission indicates molecular gas masses of 0.7 to $\\rm 2.5\\times10^{10}\\,M_{\\odot}$ in the quasar host galaxies. The observed CO line widths are spread over a wide range from $\\rm 160$ to $\\rm 860\\,km\\,s^{-1}$, with a median value of about $\\rm 360\\,km\\,s^{-1}$. The gas mass and CO line width distributions of the z$\\sim$6 quasars are consistent with samples of CO-observed SMGs and quasars at $\\rm 1.4\\leq z\\leq5$. The CO and FIR luminosities of the eight z$\\sim$6 quasars follow the $\\rm L_{FIR}-L'_{CO}$ relationship derived for local spirals, LIRGs, ULIRGs, high-z SMGs, and CO-detected quasars, though the weakest CO detection has the largest offset from the trend. This is consistent with the idea of co-eval star formation with rapid growth of the supermassive black hole in the early quasar-host systems. The derived SFRs are from $\\rm \\sim530$ to $\\rm 2380\\,M_{\\odot}\\,yr^{-1}$. The corresponding star formation efficiencies indicated by the ratios of $\\rm SFR/M_{gas}$ are consistent with the extreme starburst systems at low and high redshifts. We investigate the black hole-bulge correlations of these FIR and CO luminous quasars at z$\\sim$6 using the CO measurements. Based on certain assumptions of the molecular gas disk size, average inclination angle, and $\\sigma$-CO line width relation, we estimate the bulge dynamical masses and velocity dispersions for our sample and compare them to the local black hole-bulge relationships. The results suggest that the black hole masses of these z$\\sim$6 quasars are typically an order of magnitude higher than the values expected from the present-day relationships, which is consistent with the idea that the formation of the SMBHs occurs prior to that of the stellar bulges in the massive high-z quasar-galaxy systems. However, we also recognize that there are large uncertainties in the estimations of $\\rm M_{bulge}$ and $\\rm \\sigma$ for individual objects due to unknown gas distribution, disk inclination, and dynamics. Further high-resolution observations of quasar host galaxies should focus on these FIR and CO luminous z$\\sim$6 quasars to fully understand the black hole-bulge evolution at the highest redshift." }, "1002/1002.3614_arXiv.txt": { "abstract": "{Neglecting the second order corrections in weak lensing measurements can lead to a few percent uncertainties on cosmic shears, and becomes more important for cluster lensing mass reconstructions. Existing methods which claim to measure the reduced shears are not necessarily accurate to the second order when a point spread function (PSF) is present. We show that the method of Zhang (2008) exactly measures the reduced shears at the second order level in the presence of PSF. A simple theorem is provided for further confirming our calculation, and for judging the accuracy of any shear measurement method at the second order based on its properties at the first order. The method of Zhang (2008) is well defined mathematically. It does not require assumptions on the morphologies of galaxies and the PSF. To reach a sub-percent level accuracy, the CCD pixel size is required to be not larger than $1/3$ of the Full Width at Half Maximum (FWHM) of the PSF, regardless of whether the PSF has a power-law or exponential profile at large distances. Using a large ensemble ($\\gtrsim 10^7$) of mock galaxies of unrestricted morphologies, we study the shear recovery accuracy under different noise conditions. We find that contaminations to the shear signals from the noise of background photons can be removed in a well defined way because they are not correlated with the source shapes. The residual shear measurement errors due to background noise are consistent with zero at the sub-percent level even when the amplitude of such noise reaches about $1/10$ of the source flux within the half-light radius of the source. This limit can in principle be extended further with a larger galaxy ensemble in our simulations. On the other hand, the source Poisson noise remains to be a cause of systematic errors. For a sub-percent level accuracy, our method requires the amplitude of the source Poisson noise to be less than $1/80\\sim 1/100$ of the source flux within the half-light radius of the source, corresponding to collecting roughly $10^4$ source photons.} ", "introduction": "\\label{intro} Weak gravitational lensing has been widely used as a direct probe of the mass distribution of our Universe on different scales, including large scale structure, clusters, galaxies, etc. \\cite{hj08}. Not only is the physics of lensing well understood in the context of General Relativity, but the lensing effect can also be straightforwardly measured using the shapes of background galaxy images \\cite{kwl00,vw00,wittman00}. Currently, one of the main challenges in this field is about how to accurately recover the cosmic shear field from galaxy shapes \\cite{heymans06,massey07,bridle09a,bridle09b}. This is difficult due to the large galaxy shape noise, the involvement of the point spread function (PSF), the presence of the photon noise, the pixelation effect, etc.. There have been many literatures focusing on this particular topic \\cite{tyson90,bonnet95,kaiser95,luppino97,hoekstra98,rhodes00,kaiser00,bridle01,bernstein02,refregierbacon03,massey05,kuijken06,miller07,nakajima07,kitching08,zhang08,zhang09}. In all of the practical shear measurement methods proposed so far, there is a common assumption: the cosmic shear is small, therefore the second or higher order terms in shear can be neglected. This is true for the shear field of our Universe on large scales, which is typically of order a few percent. However, future weak lensing survey may require shear measurement accuracy to be controlled below a $0.1\\%$ level \\cite{htbj06,ar08}. On arc minute angular scales, the second order terms can cause a systematic error of order $10\\%$ to the cosmic shear power spectrum \\cite{dsw06,shapiro09}. More importantly, the shear field by a foreground cluster can easily be of order ten percent or more. Neglecting second order terms in shear measurements can lead to significant errors on the implied cluster masses \\cite{wittman01,hoekstra01,gray02,taylor04,broadhurst05,leonard07,heymans08,deb08}. The main purpose of this paper is to further develop the shear measurement method of \\cite{zhang08}(Z08 hereafter) by including the second order terms in shear/convergence in the formalism. This is done straightforwardly in \\S\\ref{method}, in which we show that to the second order in accuracy, Z08 measures exactly the reduced shears. A popular misunderstanding in the weak lensing community is that all of the existing shear measurement methods are already accurate to the second order in shear/convergence, because they all claim to measure the reduced shears. We show why this is not generally true in the presence of PSF. In \\S\\ref{theorem}, we provide a simple and useful theorem for judging whether shear estimators are accurate to the second order based on their properties at the first order. The theorem provides an easy way to see why the method of Z08 yields exactly the reduced shear. \\S\\ref{test} demonstrates the accuracy of Z08 using a large number of computer-generated mock galaxies of unrestricted morphologies in the presence of PSF and the photon noise, including both the background noise and the source Poisson noise. Finally, we conclude in \\S\\ref{summary}. ", "conclusions": "\\label{summary} Based on \\cite{zhang08,zhang09,zhang10}, we have established a robust way of measuring the cosmic shear to the second order in accuracy. The method is well defined regardless of the morphologies of the galaxies and the PSF. We have also provided a useful theorem for judging the accuracy of any shear measurement method at the second order based on its properties at the first order. For our method to achieve the accuracy at sub-percent level, the CCD pixel size is required to be not larger than about $1/3$ of the FWHM of the PSF, regardless of whether the PSF has a power-law or exponential profile at large distances\\footnote{For PSFs with strong diffraction spikes, we need to further test the method. This will be done in a future work.}. Using more than $10^7$ mock galaxies of unrestricted morphologies, we have tested the accuracy of this method under different noise conditions. We find that it is useful to separately discuss the background and source noise for any given SNR. The background noise, which is uncorrelated with the source flux, can be removed in a simple and clean way using the method of \\cite{zhang09}. In our simulations with only background noise, the shear measurement errors are found to be less than $1\\%$ for SNR as low as $10$, and the conclusion can likely be extended to even smaller SNR with simulations of a larger galaxy ensemble. On the other hand, the source Poisson noise, which strongly couples with the distribution of the source flux, remains to be the main cause of the shear measurement errors in our method. For a sub-percent level accuracy, we require the SNR of the source Poisson noise to be $\\gtrsim 80-100$. This corresponds to collecting about $10^4$ source photons per galaxy. The treatment of source Poisson noise is unclear at present, and will hopefully be addressed in a future work." }, "1002/1002.2338_arXiv.txt": { "abstract": "{Time-distance helioseismology is a technique for measuring the time for waves to travel from one point on the solar surface to another. These wave travel times are affected by advection by subsurface flows. Inferences of plasma flows based on observed travel times depend critically on the ability to accurately model the effects of subsurface flows on time-distance measurements. We present a Born-approximation based computation of the sensitivity of time-distance travel times to weak, steady, inhomogeneous subsurface flows. Three sensitivity functions are obtained, one for each component of the 3D vector flow. We show that the depth sensitivity of travel times to horizontally uniform flows is given approximately by the kinetic energy density of the oscillation modes which contribute to the travel times. For flows with strong depth dependence, the Born approximation can give substantially different results than the ray approximation. } ", "introduction": "\\sloppy Time-distance helioseismology (\\cite{Duvall1993}) is a technique for measuring the time for waves to travel from one point on the solar surface to another. Subsurface flows advect waves and as a result alter the observed travel times. Thus wave travel times can be used as probes of subsurface flows (e.g.\\ \\cite{Kosovichev1997,Zhao2004}) One of the key steps in the interpretation of travel times is to estimate the effect of subsurface flows on travel times. The ray approximation (\\cite{Kosovichev1997}), in which travel-time shifts are only caused by inhomogeneities located along the ray path connecting the observation points, has been employed in many time-distance studies of plasma flows (e.g.\\ \\cite{Kosovichev1997, Zhao2001, Zhao2003,Zhao2004}). An alternative to the ray approximation is the first Born approximation (e.g.\\ \\cite{Birch2000}; \\cite{Gizon2002}, in the context of time-distance helioseimology). The first Born approximation takes into account a single scattering and thus flows located away from the ray path can affect the travel time. \\cite{Birch2004b} studied the ranges of validity of the Born and ray approximations in a toy problem consisting of jets in a homogeneous two-dimensional medium. This study showed that there are flow configurations for which the Born approximation is valid while the ray approximation is not, especially when the transverse size of the jet is much smaller than the wavelength. As a result, we would like to use the first Born approximation to compute the sensitivity of travel times to flows for use in time-distance helioseismology. Here we employ the Born approximation approach of Gizon \\& Birch (2002), referred to as GB02 hereafter, to compute the sensitivity of travel times to weak, steady, three-dimensional subsurface flows in the Sun. We use the phenomenological model of Birch et al.\\ (2004), referred to as B04 hereafter, to describe how waves are excited and damped by convection. We work in Cartesian geometry, which is appropriate for waves that travel distances much less than the solar radius and also have wavelengths much smaller than the solar radius. The remainder of this paper is organized as follows. In \\S\\ref{sec.kernels} we derive the general expression for the sensitivity of travel times to weak flows. In \\S\\ref{sec.examples} we show a few example calculations. We compare the Born and ray approximations in \\S\\ref{sec.compare}. We conclude in \\S\\ref{sec.discuss} with a summary and a short discussion of the implications of the work presented here. ", "conclusions": "\\label{sec.discuss} We employed the Born approximation to obtain the three-dimensional sensitivity of time-distance measurements to advection by local flows. In this paper we addressed the important question of time-independent mass flows. We have shown that for horizontally uniform flows the depth dependence of the sensitivity of travel times is given approximately by the kinetic energy density of the mode which contributes most to the travel times. For simple cylindrically symmetric models of supergranulation-scale convection cells we showed that the Born and ray approximations can give results that are substantially different when the flow varies in the depth range just below the lower turning point of the ray. This suggests that for inversions of supergranulation-scale flows it may be important to use kernels based on the Born approximation rather than the ray approximation. In the future, a number of improvements could be implemented: inclusion of modes above the acoustic cutoff frequency, taking spherical geometry into account, and treatment of time-dependent flows." }, "1002/1002.0421_arXiv.txt": { "abstract": "The specific angular momentum of a Kerr black hole must not be larger than its mass. The observational confirmation of this bound which we call a Kerr bound directly suggests the existence of a black hole. In order to investigate observational testability of this bound by using the X-ray energy spectrum of black hole candidates, we calculate energy spectra for a super-spinning object (or a naked singularity) which is described by a Kerr metric but whose specific angular momentum is larger than its mass, and then compare the spectra of this object with those of a black hole. We assume an optically thick and geometrically thin disc around the super-spinning object and calculate its thermal energy spectrum seen by a distant observer by solving general relativistic radiative transfer equations including usual special and general relativistic effects such as Doppler boosting, gravitational redshift, light bending and frame-dragging. Surprisingly, for a given black hole, we can always find its super-spinning counterpart with its spin $a_*$ in the range $5/31$. In this paper, we call this object a super-spinar \\cite{gh09,bfh09}. Next, the energy spectrum of the assumed object is calculated and finally compare the spectrum with that of a black hole. The X-ray spectrum of a black hole candidate generally consists of a thermal component originating from the accretion disc around the black hole and a non-thermal component originating from high energy photons which are up-scattered in e.g. corona above or in the accretion disc. In this study, we only assume a thermal component for simplicity. The observed energy spectrum is calculated by solving the general relativistic radiative transfer including usual special and general relativistic effects such as Doppler boosting, gravitational redshift, light bending and frame-dragging. In this paper, we assume no emission from a central object. The present paper is organized as follows. In \\S\\ref{sec:disc}, physical assumptions and disc structure are given. In \\S\\ref{sec:spec}, we calculate the local radiation flux, the radial temperature profile and the energy spectrum of the disc. We give discussion in \\S\\ref{sec:dis} and conclusions are presented in \\S\\ref{sec:con}. Throughout this paper, we use the geometrical units $c=G=1$. ", "conclusions": "\\label{sec:con} The observational confirmation of the Kerr bound directly suggests the existence of a black hole. In this study, in order to investigate testability of this bound by using the observed X-ray energy spectrum of black hole candidates, we first calculate the energy spectrum for the object whose spacetime geometry is described by the Kerr metric but whose specific angular momentum is larger than its mass, and then compare the results with that of a black hole. We call this object a super-spinar in this study. The optically thick and geometrically thin disc is assumed and only the thermal energy spectrum seen by the distant observer is calculated by general relativistic radiative transfer calculations including usual special and general relativistic effects such as Doppler boosting, gravitational redshift, light bending and frame-dragging. After calculating a disc structure such as velocity fields (Fig\\ref{fig:superKerr_ELOm}) and radiation efficiency at ISCO (Fig\\ref{fig:efficiency}), we have calculated energy flux radiated from the disc (Fig \\ref{fig:PTflux}), disc temperature (Fig{\\ref{fig:T}}) and observed energy spectrum (Fig \\ref{fig:SED}). We use the new analytic formula for the radiation flux of a disc. As known in past studies, some energy is extracted from the central objects whose specific angular momentum is larger than its mass. We have investigated the influence of the extracted energy on the energy spectrum of a disc. Finally, we compare the energy spectra of a super-spinar and that of a black hole. In terms of the energy spectrum observed by a distant observer, we have obtained the following results: \\begin{itemize} \\item For the super-spinar with $18\\sqrt{6}/3 \\simeq 6.532$, the total radiation energy of the disc is lower than the disc around the non-rotating black hole. \\end{itemize} As a result of this study, we found, surprisingly, that for a given black hole we can always find its super-spinning counterpart whose observed spectrum is very similar to and practically indistinguishable from that of the black hole. Then, in order to confirm the Kerr bound we need more than the X-ray thermal spectrum of the black hole candidates." }, "1002/1002.2754_arXiv.txt": { "abstract": "We investigate the one-dimensional interaction of a relativistic jet and an external medium. Relativistic magnetohydrodynamic simulations show an anomalous boost of the jet fluid in the boundary layer, as previously reported. We describe the boost mechanism using an ideal relativistic fluid and magnetohydrodynamic theory. The kinetic model is also examined for further understanding. Simple scaling laws for the maximum Lorentz factor are derived, and verified by the simulations. ", "introduction": "Relativistic jets are considered in various contexts in high-energy astrophysics, such as active galactic nuclei (AGNs) \\citep{up95,ferrari98}, microquasars \\citep{mirabel99}, and potentially gamma-ray bursts (GRBs) \\citep{piran05,mes06}. The interaction between fast moving jets (the relevant Lorentz factors are $\\gamma_{jet} \\sim 10$--$20$ in AGNs and $\\gamma_{jet} \\gtrsim 10^2$ in GRBs) and the surrounding medium is very important to understand global dynamics of the jet system, because it is related to the mass, momentum, and energy transport across the boundary layers. In this context, development of velocity shear instabilities has been of interest (\\citet{tur76,bp76,ferrari80,birk91,bodo04,osm08} and references therein). Moreover, a relativistic jet-medium boundary is a potential site of high energy particle acceleration as well \\citep{ost00,so02}. Recently, it has been reported that the jet-medium interaction is more complex than thought even in the simplest one-dimensional (1D) case. Raising a Riemann problem of relativistic hydrodynamics (RHD), \\citet{aloy06} showed that the tangential hydrodynamic velocity and the relevant Lorentz factor ($\\gamma_{BL}$) in the boundary layer are anomalously accelerated ($\\gamma_{BL}>\\gamma_{jet}$) when the jet is over-pressured. \\citet{mizuno08} studied relativistic magnetohydrodynamic (RMHD) effect, and reported that the perpendicular magnetic field enhances the boost effect. \\citet{kom09b} discussed a similar tangential boost in their RMHD simulation of the collapser jet. Such anomalous boost effect may be responsible for increasing the jet's Lorentz factor \\citep{aloy06,mizuno08} and for modulating the radiative signature of the jet \\citep{aloy08}. However, its physical mechanism remains unclear, and therefore no quantitative analysis has been performed. In this paper, we study the mechanism of the anomalous boost by using RHD/RMHD simulations and an analytic theory. In Section 2, we describe the problem setup. In Section 3, we present the simulation results. In Section 4, we construct an RHD/RMHD theory of the problem. In Section 5, we additionally discuss kinetic aspects. The last section Section 6 contains discussions and summary. ", "conclusions": "As shown in Equation \\ref{eq:dpdt}, the anomalous bulk boost comes from the temporal decrease of relativistic pressure. From the energy viewpoint, the term transports the internal energy to that of the bulk motion ($p_g \\Rightarrow \\gamma$), as mentioned by \\citet{aloy08}. The internal-to-bulk energy transport is somewhat counter-intuitive, however, it is a logical consequence of the relativistic fluid formalism. The site of the boost is the rarefaction region. The rarefaction wave involves the temporal pressure decrease behind its wave front and there is a room for the convective fluid motion (Equation \\ref{eq:p_force}). In contrast, neither conditions are satisfied around the shocks. The anomalous boost does not occur on the other side of the contact/tangential discontinuity nor will it occur when another shock replaces the rarefaction wave. Therefore, the transition from the shock regime to the rarefaction wave regime \\citep{rez02} would be a critical condition for the problem. Similar boost in the normal direction has recently been reported in magnetically-dominated rarefaction region as well \\citep{mizuno09}. Another explanation is a relativistic free expansion in the jet frame \\citep{kom09b}. When the relativistically strong pressure pushes the gas outward against the external medium, the lateral expansion can be relativistic in the jet frame. Then, the Lorentz factor in the observer frame yields $\\gamma_{BL} \\sim \\gamma_{jet} (1-v'^2)^{-1/2}$, where $v'$ is the expansion speed in the jet frame. We expect that the term $-\\vec{v}'\\partial_{t'} p_t$ enhances such expansion in the rarefaction region, and that the relevant boost is projected into the tangential boost in the observer frame. Strictly speaking, a 1D problem in the observer frame is no longer identical to that in the jet frame, because a 1D expansion of the discontinuity front in the +$x$-direction is projected to the oblique direction in the jet frame. The two problems start differently and therefore the situation is more complicated. A potential limitation is that multi-dimensional instabilities may modulate the 1D evolution. Especially, the relativistic Kelvin--Helmholtz (KH) instabilities will be relevant. In the regime of our interest, the increasing Lorentz factor \\citep{tur76,bp76,bodo04} and the flow-aligned magnetic field \\citep{osm08} suppress the KH mode; for instance, if we employ \\citet{bodo04}'s stability condition of $\\gamma_{jet} > ( 1+2 \\cos^{-2}{\\theta} )$ in our RHD jet ($\\gamma_{jet}=7$), where $\\theta$ is the angle between the jet flow and the wavevector, the instability is allowed only in the quasi-transverse direction. On the other hand, shear layers with density asymmetry are known to be substantially KH-unstable. Once the KH vortex develops, the subsequent turbulence is likely to smooth the sharp lateral structure. While 1D-like signatures have been found in some three-dimensional RHD \\citep{aloy05} and two-dimensional RMHD simulations \\citep{mizuno08,tch09,kom09b}, interference with the KH and other instabilities needs further investigation. In addition, we need to keep in mind that the entire process depends on the ideal fluid assumption. In order to justify it, collisional or other scattering processes have to relax the gas much quicker than the dynamical timescale. However, those are difficult conditions especially in the jet side, where the physical processes look even slower by the relativistic effect. In Section 5, we show that the fluid bulk speed is considerably smaller than a wide thermal spread in the momentum profile, when the boost operates. We think that the counter-intuitive force may be just enforced by the ideal fluid assumption: i.e. the anomalous fluid acceleration may be an artifact of an expedient isotropic fluid velocity. In the real world, we expect that non-ideal effects such as the heat flow play roles. In fact, the system involves large gradient of the pressure and the temperature in the rarefaction regions and around the discontinuities. In the high-temperature regime of $T \\gg 1$, the energy and momentum balances are mainly controlled by the pressure parts (the internal energy or the enthalpy flux), which can be sensitive to the local gas distribution functions. In summary, we examined the 1D anomalous relativistic boost \\citep{aloy06,mizuno08} at the lateral boundary of relativistic jets. We numerically and theoretically confirmed that the anomalous boost occurs in the RHD and RMHD regimes. We further derived simple scaling laws for the accelerated Lorentz factor, \\begin{eqnarray} \\gamma_{BL} \\lesssim \\gamma_{jet} \\Big(\\frac{p_{t,L}}{p_{t,R}}\\Big)^s \\left\\{ \\begin{array}{cl} s=1/4 & ~~({\\rm hydro,~parallel})\\\\ s=1/2 & ~~({\\rm perpendicular}) \\\\ \\end{array} \\right. \\nonumber \\end{eqnarray} We also note that the process operates in an {\\itshape ideal} fluid. The non-ideal effects (heat flow etc.) as well as multi-dimensional effects are left for future works. We hope that this work will be a basic piece for the boundary problems in relativistic jets and the relevant simulations." }, "1002/1002.0506.txt": { "abstract": "{Simulations of astrophysical disks in the shearing box that are subject to the magnetorotational instability (MRI) show that activity appears to be reduced as the magnetic Prandtl number P$_{{\\rm m}}$ is lowered. It is therefore important to understand the reasons for this trend, especially if this trend is shown to continue when higher resolution calculations are performed in the near future. Calculations for laboratory experiments show that saturation is achieved through modification of the background shear for P$_{{\\rm m}} \\ll 1$.} {Guided by the results of calculations appropriate for laboratory experiments when P$_{{\\rm m}}$ is very low, the stability of {inviscid} disturbances in a shearing box model immersed in a constant vertical background magnetic field is considered under a variety of shear profiles and boundary conditions in order to evaluate the hypothesis that modifications of the shear bring about saturation of the instability. Shear profiles $q$ are given by the local background Keplerian mean, $q_0$, plus time-independent departures, $Q(x)$, with zero average on a given scale.} {The axisymmetric linear stability of {inviscid} magnetohydrodynamic normal modes in the shearing box is analyzed.} {(i) The stability/instability of modes subject to modified shear profiles {may be interpreted} by a generalized Velikhov criterion given by an effective shear and radial wavenumber that are defined by the radial structure of the mode and the form of $Q$. (ii) Where channel modes occur, comparisons against marginally unstable disturbance in the classical case, $Q=0$, shows that all modifications of the shear examined here enhance mode instability. (iii) For models with boundary conditions mimicing laboratory experiments, modified shear profiles exist that stabilize a marginally unstable MRI for $Q=0$. (iv) Localized normal modes on domains of infinite radial extent characterized by either single defects or symmetric top-hat profiles for $Q$ are also investigated. If the regions of modified shear are less (greater) than the local Keplerian background, then there are (are no) normal modes leading to the MRI.} {The emergence and stability of the MRI is sensitive to the boundary conditions adopted. Channel modes do not appear to be stabilized through modifications of the background shear whose average remains Keplerian. However, systems that have non-penetrative boundaries can saturate the MRI through modification of the background shear. Conceptually equating the qualitative results from laboratory experiments to the conditions in a disk may therefore be misleading.} % context heading (optional) % {} leave it empty if necessary % {Nonsense} % {Nonsense} % {Nonsense} % { Nonsense} % {Nonsense} \\titlerunning{Speculation on low magnetic Prandtl number disk velocity profiles} ", "introduction": " \\begin{enumerate} \\item Numerical experiments of low magnetic Prandtl number flows appropriate for cold disk environments are outside current computational reach. \\item Current numerical experiments seem to show the trend that as P$_{{\\rm m}}$ is lowered so is the corresponding angular momentum transport. The reasons for this is not yet clear. \\item Theoretical analysis of laboratory setups evaluating the response of the MRI subject to a vertical background field find that transport is reduced with P$_{{\\rm m}}$. Furthermore this reduction emerges through the modification of the basic shear profile (e.g. Taylor-Couette) inside the experimental cavity which, in turn, is driven into place by the MRI itself. The final condition is a stable pattern state. \\item In the final pattern state reached by the model laboratory system the modified shear profile's amplitude is independent of P$_{{\\rm m}}$. This is not true of the remaining fluid and magnetic quantities which scale as P$_{{\\rm m}}^{\\lambda}$ with typical values $\\lambda \\approx 1/2$. As P$_{{\\rm m}}$ {is made very small one would observe a fluid state characterized by a uniform background magnetic field and an azimuthal flow showing deviations from a Keplerian state.} \\item The main mode of instability and driver of turbulence in the numerical experiments of the shearing box is the channel mode which is a special response with no radial structure. This mode precipitates a secondary instability which is understood to lead to a turbulent cascade. There is no channel mode allowed in a setup appropriate for a laboratory experiment. \\end{enumerate} The saturation hypothesis is then the following: is it possible that the fluid response in a shearing sheet environment characterized by very low P$_{{\\rm m}}$ stabilizes itself by modifying the background shear as suggested by the calculations appropriate for laboratory conditions? Of course, there are some constraints placed upon the way in which the shear may be modified in that the gravitational field predominantly ``forces\" there to be a Keplerian profile (Ebrahimi et al. 2009). However by analogy to the situation for the laboratory setup, there may come into play undulations of the shear whose average is zero over some length scale $L$ (e.g. Figure 1). \\par Of course, the main difference between the laboratory setup and numerical calculations of the shearing box are their respective boundary conditions. The numerical experiments adopt periodic boundary conditions while the calculations for the laboratory experiments require no normal flow conditions at an inner and outer radius and the former of these permits channel modes to come into play. \\par Thus, what has been done here is a simple model calculation in which a shearing box configuration with a constant background vertical magnetic field is tested for stability for a number of different shear profiles. The perspective taken is that if there exists a shear profile (say, whose average over some length scale is zero) that is stable to the MRI, then it would strongly suggest that low P$_{{\\rm m}}$ flow configurations might saturate the MRI by driving into place a new azimuthal flow profile which leads to saturation. \\par What is found here is something very interesting. For flow configurations in which no-normal flow conditions are imposed at both radial boundaries (Sections \\ref{ChannelWalls_smallQ} \\& \\ref{finite_domain_single_defect_channel_walls}) there exists a modification to the shear in which the most unstable MRI mode is stabilized. However, for flows in which periodic boundary conditions (Sections \\ref{small_Q_periodic_conditions} \\& \\ref{ShearDefectPeriodicDomain}) or fixed Lagrangian pressure conditions (Section \\ref{small_Q_pressure_conditions}) are imposed there are no modulations of the shear which shut off the instability. In these cases, all modifications of the shear (with zero mean on some radial length scale) seem to enhance instability of the channel mode. For example, if the given channel mode is marginally stable to exponential growth, then introduction of the modified shear, no matter what is its non-zero amplitude, seems to destabilize it as the results of Section \\ref{small_Q_periodic_conditions} and Section \\ref{ShearDefectPeriodicDomain} indicate.\\par This strongly implies that in those numerical experiments in which a shear is threaded by a constant vertical magnetic field, the shear modification/MRI-stabilization process may only apply for those flow conditions in which there is no normal flow on either one or both radial boundaries of the system. \\par On the other hand, for a similar magnetized fluid configuration described by periodic boundaries which therefore permits channel modes to exist, modifications of the shear appears to further destabilize the channel mode. If one were to consider a flow configuration (e.g. the local Keplerian profile in a shearing box threaded by a constant background field) that admits an unstable channel mode, then there seems to be no modification of the shear with zero mean that would saturate the growth of the channel mode \\emph{since all such modifications further enhances its destabilizing influence}. The emergence of a pattern state leading to saturation like the one envisioned by the hypothesis would seem to be out of the question under these circumstances. Moreover this suggests that the tendency for angular momentum transport reduction with P$_{{\\rm m}}$, as indicated by numerical experiments of the shearing box, is due to some other process or processes (which are not discussed here). \\par These issues deserve further investigation as the conclusions reached here were done so utilizing very simple functional forms for the shear profile $Q$. Namely, the analysis performed in the previous sections either assumed weak values of $Q$, in order to use singular perturbation theory techniques, or that $Q$ is described by delta-functions (see also below). Use of these functions, which greatly facilitate analysis, have offered some insights into the nature of the responses of these modes. Of course, further investigation would require considering more physically realizable profiles and to check the robustness of the results they yield against the ones offered by use of these more simplified forms. As such the results reached here can serve as guideposts for further enquiries along these lines. \\par In the following some reflections are presented on other issues pertaining to the implications of the calculations performed here. \\subsection{On the use of delta-functions} A comparison between the responses of a single delta-function shear profile on an infinite domain (Section \\ref{single_defect}) and a top-hat profile for $Q$ (Section \\ref{symmetric_shear_step}) show that the results of the two configurations are qualitatively similar. Although the symmetric step function allows for more normal modes to exist as its radial extent is increased - a feature which is missed if one uses the delta-function prescription - the qualitative flavor of the response is captured by its use nevertheless. \\par Another deficiency found in using the defect prescription is that for the top-hat profile studied in Section \\ref{symmetric_shear_step} there are both even and odd parity normal modes possible whereas only the even-parity mode is admitted in the defect model. The odd-parity mode \\emph{has no counterpart} in the single defect model. It was also shown that the odd-parity mode may exist if the radial domain of the symmetric shear step is large enough as (\\ref{min_L_criterion_odd_parity_mode}) indicates that the minimum size of the domain required for this mode to exist has the proportionality $L\\sim1/\\sqrt{Q_0}$. For small values of the defect amplitude the minimum size of the shear step required for this odd-parity mode to exist is so large that one does not expect that this should play any physically relevant role in the dynamics if one is interested in the response to narrow regions of altered shear. It is with some confidence to suppose that it is reasonable to represent the dynamical response of slender regions of modified shear through the use of delta-function defects and this is the motivation and justification for its use in Sections 5.5 and 5.6. \\subsection{Channel modes} In the small $Q$ theory of Section \\ref{small_Q_theory}, it was shown that channel modes exist for periodic boundary conditions and for conditions in which the total pressure perturbations are zero at the two boundaries. Conversely, if either one of the two boundaries force the velocity perturbation to be zero there, then there is no channel mode permitted as a normal mode solution. In light of the previous results it is important that the boundary conditions appropriate for disks be properly ascertained in order to assess and better understand how activity in them are driven.\\par Channel modes in the classical limit have no radial structure. In Section \\ref{ShearDefectPeriodicDomain} the development of this mode was studied as a function of the defect amplitude $Q_0$ and it is shown that though the channel mode develops some amount of radial structure, as evinced by the non-zero value of its wavenumber $\\kappa_{_{00}}$, its tendency to be unstable survives and is even enhanced. In Regev \\& Umurhan (2008) it was postulated that the channel mode might disappear as a natural mode of the system if some amount of radial symmetry breaking were introduced into the governing equations of motion, for instance, in the form of a radially varying shear profile. However, the results of this section show that the classical channel mode indeed persists in its existence despite the introduction of a shear profile which deviates from the background constant shear state with zero average over some length scale. These conclusions are consistent with similar results regarding the nature of the MRI in more global contexts (Curry et al. 1994). \\subsection{On the existence and number of localized normal modes} In Sections \\ref{single_defect} and \\ref{symmetric_shear_step} the existence of localized modes was examined. Localization is understood in this case to correspond to modal disturbances which show exponential decay as the radial coordinate $x\\rightarrow \\pm\\infty$. Thus, by construction, incoming/outgoing waves are excluded from consideration. \\footnote{Given the mathematical structure of the ode describing axisymmetric disturbances, the radial eigenfunctions that are solutions have either strictly oscillatory or exponential radial profiles. Thus the possibility of decaying oscillations is ruled out.} The existence of localized normal modes depends on the relative sign of the shear defect/step profile $Q_0$: HMI modes exist if the shear in the defect/top-hat region is stronger than the background ($Q_0 > 0$) while HI modes exist if the shear is relatively weakened ($Q_0 < 0$) \\footnote{The mode existence dependency upon the sign of $Q_0$ was checked for other shear profiles (for instance, by replacing the top-hat with a truncated parabolic profile). Aside from differences in the magnitude of the temporal response and the multiplicity of modes allowed, the aformentioned sign dependency still holds.} - summarized here as \\beqa & & Q_0 >0 \\longleftrightarrow {\\rm hydromagnetic} \\ {\\rm inertial} \\ {\\rm disturbances}, \\nonumber \\\\ & & Q_0 <0 \\longleftrightarrow {\\rm hydro-inertial} \\ {\\rm disturbances}. \\nonumber \\eeqa Given that the former mode type leads to the MRI, it means that if the shear in the defect zone is weakened, then there are no localized normal modes which can go unstable as only hydro-inertial modes are supported (which happen to be stable except for very extreme circumstances). Of course, as an initial value problem, it does not mean that there are no hydromagnetic inertial modes permitted but, instead, the mode probably has some type of algebraic temporal growth/decay associated with it which cannot be described by the usual normal mode approach. This needs to be further examined by studying the associated initial-value problem. Similar features are known to exist for vertically localized MRI modes in disks (Liverts \\& Mond, 2009) in which the response of the initial value problem can exhibit exponential growth with an amplitude supporting some algebraic time-dependence. Growth can initially be very small and it can take a very long time for a given disturbance to reach the same amplitude as expected for its counterpart normal mode. A conjecture is that the meaning of the absence of normal modes in the problem considered here may have similar attributes to the features associated with the initial-value problem examined by Liverts \\& Mond (2009). \\par Irrespective of this outcome it should be remembered that the dichotomy observed here for modes on an infinite radial domain is a consequence of the imposition of exponential decay of the disturbances. This distinction disappears once waves are allowed to enter and exit from the ``infinite\" boundaries.\\par {Lastly it is also noteworthy that for the localized problems considered here, when normal modes are permitted they usually come as a finite set. This is in contrast to other localized flow problems considered elsewhere in the literature. For example, in the problem of compact rotating magnetized jets in an infinite medium considered by Bodo et al. (1989) a countably infinite number of normal-modes are predicted for that system. This has some origin in the uniformity of both the rotation and magnetic field profiles of the jet considered in that study. By contrast, there are examples in geophysical flows in which the number of normal mode disturbances are discrete and finite. A recent example can be found in the work of Griffiths (2008) where the stability of inertial waves are studied in a stratified rotating flow with strong horizontal shear. For strongly peaked forms of the potential vorticity only a finite set of normal modes are predicted for the system. Such finiteness of the number of normal modes permitted is not so unusual in systems in which background states, like the shear or stratification, have strongly peaked functional forms like they are in the cases studied in this work.} \\subsection{Stabilization in general and an interpretive tool} The absolute minimum condition that must be met for any normal mode of the system to be unstable in a configuration with constant shear $q_0$ is given by the Velikhov criterion (\\ref{velikhov_condition}): if the square of the Alfven frequency ($\\omega_{az}^2$) is greater than the shear by $2\\Omega_0^2 q_0$, then none of the HMI modes of the system can be unstable. For individual modes in the same system with some amount of radial structure the criterion is given by (\\ref{criterion_for_instability_classic_mri}), where $\\beta_n^2$ is the square of the radial wavenumber of the disturbance. As can be seen, for fixed values of $\\omega_{az}^2$ there are two ways in which stability may be promoted - either by weakening the shear or by increasing the radial wavenumber. In the problem where the shear is constant, these quantities are set from the outset as parameters. \\par {At the end of Section 4 there is a series of arguments for arbitrary shear profiles which lead to the identity found in (\\ref{genA}), which is like a generalized eigenvalue condition. } This relationship is the analog of the dispersion condition of the classical limit (uniform $q$) contained in (\\ref{dispersion_condition_uniform_q}).\\footnote{The classical limit is obtained when the replacements $\\bar\\beta \\rightarrow \\beta_n$ and $q\\rightarrow q_0$ are made.} { Care must be adopted before interpreting this expression as an eigenvalue condition since $\\sigma$ appears implicitly in the expressions for $\\bar q$ and $\\bar\\beta$.} Nonetheless, a parallel analysis shows that stability is promoted if \\beq \\varpi \\equiv \\omega_{az}^2 - \\frac{2\\Omega_0^2 \\bar q k^2}{\\bar\\beta^2 + k^2} > 0, \\label{interpretive_tool} \\eeq where the effective wavenumber $\\bar\\beta$ and shear $\\bar q$, given in (\\ref{genB},\\ref{genC}) are, \\[ \\bar\\beta^2 \\equiv \\frac{\\int_{{\\cal D}}|\\partial_x\\psi|^2 dx} {\\int_{{\\cal D}}|\\psi|^2 dx}, \\qquad \\bar q \\equiv q_0 + \\frac{\\int_{{\\cal D}}Q|\\psi|^2 dx} {\\int_{{\\cal D}}|\\psi|^2 dx}. \\] ${\\cal D}$ is the domain and $\\psi$ is the radial eigenmode of the disturbance in question. The stability condition for general $q$ may be rationalized in the same way as the criterion is understood for the classical case with the replacements, $\\beta_n \\rightarrow \\bar\\beta$ and $ q_0 \\rightarrow \\bar q$. {With these aforementioned observations and caveats in mind, one may use (\\ref{interpretive_tool}) as an interpretive tool to aid in understanding the reasons for stability/instability in a particular problem.} Thus, stability is promoted if (i) for fixed effective wavenumber, the effective shear is reduced or (ii) for fixed effective shear the effective wavenumber is increased, or in more general terms (iii) if the combination results in an overall reduction of the quotient \\[ \\frac{\\bar q}{\\bar\\beta^2 + k^2}. \\] If one considers the results concerning the problem with periodic boundaries in Section \\ref{ShearDefectPeriodicDomain}, then the persistent instability expected for channel modes can be understood in terms of this criterion. In the classical limit where the defect amplitude is zero then $\\bar\\beta$ for the channel mode is zero. As the defect amplitude is raised $\\bar\\beta$ increases suggesting that this effect has a stabilizing influence, e.g. see Figure \\ref{section6plot1}. However, as the defect amplitude is raised away from zero the effective shear increases no matter what the sign of $Q_0$ and, to effect, this increase always sufficiently overpowers the increase in $\\bar\\beta^2$ so that the end result is further instability. In the same figure is plotted the quantity \\beq \\frac{\\varpi}{\\varpi_0} = \\frac{\\bar q}{q_0}\\frac{\\beta_n^2 + k^2}{\\bar\\beta^2 + k^2}, \\eeq which, for the generalized Velikhov criterion, is a measure of the ratio of the destabilizing term for the shear profile $q$ compared to its nominal value when the shear is only $q_0$: for ${\\varpi}/{\\varpi_0} > 1$ one expects enhanced destabilization against the classical mode while for ${\\varpi}/{\\varpi_0} < 1$ comes with enhanced stabilization. In Figure \\ref{section6plot1} this quantity is plotted for one of the channel modes discussed in \\ref{ShearDefectPeriodicDomain}. Noting that in this case $\\beta_n = 0$ (i.e. that its radial wavenumber is zero in the ideal case) the radial profile of the mode develops as the amplitude of the shear defect increases no matter what its sign. The stabilizing rise in $\\bar\\beta$ is accompanied by the destabilizing increase of the effective shear, that is to say $\\bar Q > 0$, so that the quotient of their respective influences results in ${\\varpi}/{\\varpi_0} > 1$, indicating instability. \\par Similarly plotted in Figure \\ref{section6plot2} is the influence $Q_0$ has upon $\\bar\\beta$, $\\bar q$, and the stability measure $\\varpi/\\varpi_0$, for the single defect problem on a finite domain studied in Section \\ref{finite_domain_single_defect_channel_walls}. The presence of walls filters out the channel mode and the resulting flow potentially can support shear profiles which can stabilize the least stable mode. The Figure shows that within the range of (negative) values of the defect amplitude $Q_0$ for which the mode does not exponentially grow/decay the effective shear is less than the background shear state only for part of the stable range. It means to say, then, in that range of parameter values in which the effective shear is greater than the background shear, yet there is still stability of the mode, stabilization is brought through an increased effective radial wavenumber which over-compensates the increased destabilization brought about by an increase in the effective shear. \\begin{figure} \\begin{center} \\leavevmode \\epsfysize=8.15cm \\epsfbox{13169fg9.eps} \\end{center} \\caption{The behavior of the unstable channel mode for the problem with periodic boundary conditions examined in Section \\ref{ShearDefectPeriodicDomain}: $kL = \\pi$ and $\\omega_{az}^2 = 2q_0$ with $q_0 = 3/2$. Plotted is the temporal response, $\\sigma^2$, together with its corresponding effective wavenumber and shear $\\bar \\beta$ and $\\bar q = q_0 + \\bar Q$ ($\\bar Q$ plotted), and the stability measure for this mode $\\varpi/\\varpi_0$. All values of the shear defect $Q_0$ correspond to increasing the effective shear to above the background value $\\bar Q(Q_0) > 0$. This increase in $\\bar q$ always overpowers the stabilizing influence of an increased $\\bar\\beta^2$. Note, the values of $\\sigma^2$ have been scaled by an arbitrary factor in order to facilitate clear comparison.} \\label{section6plot1} \\end{figure} \\begin{figure} \\begin{center} \\leavevmode \\epsfysize=8.15cm \\epsfbox{13169fg10.eps} \\end{center} \\caption{The behavior of the least stable HMI mode for the problem with channel walls examined in Section \\ref{finite_domain_single_defect_channel_walls}: $kL = \\pi$ and $\\omega_{az}^2 = 2q_0/(k^2 + \\beta_1^2)$ with $\\beta_1 = \\pi/L$ and $q_0 = 3/2$. Plotted is the temporal response, $\\sigma^2$, together with its corresponding effective wavenumber and shear $\\bar \\beta$ and $\\bar q = q_0 + \\bar Q$ ($\\bar Q$ plotted), and the stability measure for this mode $\\varpi/\\varpi_0$. Stability occurs for $Q_0 \\in (-1.32, 0)$. In this window the effective shear is less than the background value, ($\\bar q - q_0 = \\bar Q < 0$) for only part of this range, i.e. for $Q_0 \\in (-0.5,0)$. Stabilization in the range $Q_0 \\in (-1.32,-0.5)$ is achieved by a sufficient \\emph{increase} of the effective wavenumber so that the overall effect upon the stabilization parameter $\\varpi/\\varpi_0$ is for it to be less than one. $\\sigma^2$ has been similarly scaled as in the previous figure.} \\label{section6plot2} \\end{figure} \\subsection{Final reflections and an implication} The fundamental basis of this study is the assumption that in low magnetic Prandtl number flows for model environments like that considered in this study, the only noticeable outcome of the development of the MRI is to alter the basic shear profile. All other quantities, such as magnetic fields and the radial and vertical velocities, have saturated profiles which scale as some positive power of P$_{\\rm m}$ so that for P$_{\\rm m} \\ll 1$ their contributions become negligible. The basis of this assumed trend derives from the theoretical calculations of laboratory setups discussed in the Introduction. It is therefore assumed here that a similar trend \\emph{may} appear for configurations in shearing boxes with periodic radial boundary conditions as well. The perspective taken here with regards to those configurations is that if this is indeed the case, then one can (in principle) test the stability of a wide variety of shear profiles which show deviations (constrained to have a zero average over some length scale) from the basic Keplerian profile - configurations which are akin to those predicted from the calculations of laboratory setups. If a stability calculation shows that there exists a shear profile which is stable to these axisymmetric disturbances, then the shear-modification/mode-saturation hypothesis should become a serious candidate explanation.\\par However, the calculations done in this study indicate that for periodic shearing box environments there is always an instability of the basic channel mode no matter what shear profile is assumed (satisfying the aforementioned constraints). Unlike the situation for laboratory setups, where saturation can be achieved by altering the basic shear profile, saturation (if present) in shearing box environments with periodic boundary conditions \\emph{is likely not due to this mechanism}.\\par This further implies that the nonlinear response of systems supporting the MRI stongly depends upon the boundary conditions employed which is not so surprising as many physical systems in Nature exhibit this sensitivity to their boundaries. Thus it seems to this author that some care must be taken before one equates the results of laboratory experiments to those physically related/analogous systems for which the experiments are meant to represent (cf. Ji et al. 2006). This cautionary note is not to diminish the value of one over the other, rather, it is intended to emphasize that there are many subtle and, at times, conflicting features between such systems that must be clearly understood before any firm conclusions are reached.", "conclusions": "This study is devoted to examining the axisymmetric inviscid linear response of a shearing sheet environment threaded by a constant vertical magnetic field for an array of different shear profiles. The motivation for considering this setup" }, "1002/1002.2562_arXiv.txt": { "abstract": "The bar formation is still an open problem in modern astrophysics. In this paper we present numerical simulation performed with the aim of analyzing the growth of the bar instability inside stellar-gaseous disks, where the star formation is triggered, and a central black hole is present. The aim of this paper is to point out the impact of such a central massive black hole on the growth of the bar. We use N-body-SPH simulations of the same isolated disk-to-halo mass systems harboring black holes with different initial masses and different energy feedback on the surrounding gas. We compare the results of these simulations with the one of the same disk without black hole in its center. We make the same comparison (disk with and without black hole) for a stellar disk in a fully cosmological scenario. A stellar bar, lasting 10 Gyrs, is present in all our simulations.\\\\ The central black hole mass has in general a mild effect on the ellipticity of the bar but it is never able to destroy it. The black holes grow in different way according their initial mass and their feedback efficiency, the final values of the velocity dispersions and of the black hole masses are near to the phenomenological constraints. ", "introduction": "In a series of recent papers \\citet{Cu06}, \\citet{Cu07}, \\citet{Cu08} we have shown that the bar formation inside disk galaxies is triggered also by the cosmological scenario and not only by the classical instability of a self gravitating disk. We obtained such result by analyzing the growth of a bar instability in exponential disks, embedded in a Dark Matter (DM) halo which evolves in a cosmological context. We used purely stellar disks as well as disks containing both gas and stars. In the latter case, we studied the effect of only allowing the gas to radiatively cool, and we then examined the impact of the star formation. Our embedding technique allowed us to vary the baryon to DM mass ratio onto the {\\it same} DM halo. In almost all the cases we studied, including the classically stable ones, we noticed the formation of a long--lasting bar. The onset of the bar instability, in the latter cases, appears to be driven by the triaxiality and by the dynamical evolution of the DM halo. For a few cases, cooling gas proved to be able to stabilize the gas, but such effect disappeared when star formation was turned on, because gas converts in stars before a stabilizing, dense, central knot can form.Therefore the problem of inhibiting the bar formation is still an open astrophysical problem, since one third at least of disk galaxies are not barred. The observational data indicate that massive central black holes exist in disk galaxies as well as ellipticals. Moreover black holes masses are correlated with host galaxies properties; \\citet{Mag98} has concluded that the median black holes mass is 0.006 of the bulge mass and \\citet{Kor01} correlated the black hole mass to the bulge velocity dispersion. Such a large mass concentration inside the galaxy could affect its structure. It is known indeed that a central mass concentrations (CMC) can destroy a bar (see e. g. \\citet{Bou05}, \\citet{Cu07} ). The effect of the presence of a massive black hole (BH) at the center of a disk galaxy could therefore mimic the stabilizing action of a knot of dense, cold gas which forms if gas is allowed to radiatively cool but not to form stars. The influence of central BHs on the dynamical evolution of bars in disk galaxies has been examined mainly by \\citet{she04} and \\citet{Hoz05} in a pure non dissipative scenario and using as model galaxies isolated systems. They results give very high values ($10^{8.5}$) the minimum mass necessary for bar dissolution, higher than the inferred one for local spirals. In this paper we want to treat the same problem using simulations of stellar- gaseous disk with star formation. We will evolve the disk model as isolated system and in a cosmological scenario. We will discuss the evolution of the bar ellipticity and semi-major axis of the same disk endowed with different initial BH mass. The BH is accreting during the evolution. Its accretion is regulated by a suitable feedback parameter(\\citet{DiMat03}), which states a coupling between a percentage of the accretion radiating energy and the gas heating. We will also show the role of this feedback parameter and its interplay with the initial BH mass, the formation of the CMC and with the star formation. Finally, we will investigate the relation between the decrease of the bar ellipticity,the accretion time and the feedback. Moreover we will compare our simulations with some phenomenological issues as the Kormendy black hole mass-bulge velocity dispersion relation. The plan of the paper is the following. In Section 2, we summarize our recipe for the initial $disk+halo$ system and present the star formation recipe. In Section 3, we present our simulations, and in Section 4 we point out our results. The parameters related to the bars formed in the new and the old stars and to the global, old+new, stellar populations are given in this section. In section 5 there are phenomenological implications and section 6 is devoted to our discussion and conclusions. ", "conclusions": "We have presented 11 isolated simulations and 2 cosmological simulations of the evolution of a stellar+gaseous disk embedded in a DM halo,with the same disk--to--halo mass ratios. The aim was to evaluate the impact of a massive black hole, with different accretion histories on the formation and evolution of a stellar bar. The old star component shows a long-lasting bar, 10 Gyrs old, in all our simulations, regardless of the presence, the mass and the energy feedback efficiency of a central BH. We noticed a mild impact of the BHs having masses greater than $10^8 h^{-1} M_{\\odot}$ on the old star component ellipticity, impact which appear stronger on the bar ellipticity of the new stars. In a previous paper (\\citet{Cu08}) we found that the star formation, by reducing the central gaseous mass concentration, allows the bar to survive until the end of the evolution in massive disks, at variance with the results in \\citet{Cu07} ( where the gas was not allowed to form stars): in the latter case a gas fraction 0.2 was able to destroy the bar. In this work we show that a BH placed and growing at the center of the disk has a small influence on the ellipticity of the bar developing in the preexistent star population. On the other hand, the BH with its feedback has an action on the formation of the CMC, which has a major role in quenching the bar. The bar formed by the newly formed stars is always weaker than in the disk without BH, independently on the BH's evolution and feedback. Also in this case, however, the presence of the BH does not completely stop the bar formation. A direct comparison between the results of our $N$-body simulations and the results of \\citet{Hoz05} is not possible. The reason is that their numerical model consists in a razor-thin disk model, without bulge and halo. It is evolved using a self-consistent field method (\\citet{Hern92}), and not with a $N$-body code. Whereas the Hozumi's disk is stabilized by a black hole having mass of $ 10^{8.5} M_{\\odot}$ our disk forms a bar even if the central BH has a higher mass. Adding gas and star formation to the same disk, we notice that an accreting massive BH has an impact on the bar ellipticity but it is not able to destroy it completely. The effect of the BH is more evident on the {\\it inner} bar ellipticity, which is reduced also when the effect of a star-forming gaseous component is considered. A comparison with the phenomenology confirms that the default feedback parameter value 0.05 seems the more appropriate to reproduce the correlations between the velocity dispersion and the black holes masses observed by Kormendy. {\\bf Acknowledgments} Simulations were performed on the CINECA IBM CLX cluster, thanks to the INAF-CINECA grants cnato43a/inato003 `` Simulations of disk galaxies in a cosmological framework: the impact of the central Black Hole''. We wish to thank V. Springel for kindly providing us with his code GADGET. We thank an anonymous referee for his/her valuable suggestions, helpful to improve the paper." }, "1002/1002.3172_arXiv.txt": { "abstract": "We explore the use of simple star-formation rate (SFR) indicators (such as may be used in high-redshift galaxy surveys) in the local Universe using \\oii, \\ha, and $u$-band luminosities from the deeper 275 deg$^2$ Stripe 82 subsample of the {\\it Sloan Digital Sky Survey} (SDSS) coupled with UV data from the {\\it Galaxy Evolution EXplorer satellite} ({\\it GALEX}). We examine the consistency of such methods using the star-formation rate density (SFRD) as a function of stellar mass in this local volume, and quantify the accuracy of corrections for dust and metallicity on the various indicators. Rest-frame $u$-band promises to be a particularly good SFR estimator for high redshift studies since it does not require a particularly large or sensitive extinction correction, yet yields results broadly consistent with more observationally expensive methods. We suggest that the \\oii-derived SFR, commonly used at higher redshifts (z$\\sim$1), can be used to reliably estimate SFRs for ensembles of galaxies, but for high mass galaxies (\\ms$\\gsim$10), a larger correction than is typically used is required to compensate for the effects of metallicity dependence and dust extinction. We provide a new empirical mass-dependent correction for the \\oii-SFR. {\\bf Note: This astro-ph version has been updated to correct the typographical errors in equations 2, 7, and 9 of the originally published paper, as described in the MNRAS Erratum.} ", "introduction": "The most useful observables in constructing a picture of galaxy formation are likely to be: a galaxy's star formation rate (SFR), its stellar mass, and the epoch at which it is observed. Stellar masses are now routinely measured out to z$\\sim$4 (e.g., \\citealt{Marchesini:2009fx}) by fitting spectral energy distributions (SEDs) of stellar population synthesis models to multicolour broadband photometry. Although the uncertainties in stellar mass estimates are dominated by evolutionary uncertainties within the models \\citep{Marchesini:2009fx}, most workers use the same models and similar passbands for the photometry (or the differences between models are at least well understood) and thus the comparison between stellar masses used in different works is relatively straightforward. The situation is somewhat less uniform for SFR measurements. Indeed, most of the systematic errors affecting SFR measurements are likely to be functions of stellar mass \\citep{Brinchmann:2004ct,Kewley:2004vs,Cowie:2008ob,Pannella:2009uj}. Over the last decade, measuring the star formation rate density (SFRD) of large samples of galaxies out to z$\\sim$6 has become routine, and dozens of such measurements now exist \\citep{Lilly:1996la,Madau:1996ao,1999ApJ...519....1S,Hopkins:2006bv,Reddy:2009wc}. In this time of `survey science', statistical samples are commonplace and the largest remaining uncertainties in such studies are systematic in nature, for example, the choice of SFR indicator. Common proxies for SFR include the luminosity of emission lines such as \\ha, or the UV continuum, which are sensitive to the ionising flux from young stars; or mid- or far-infrared radiation (MIR, FIR) which measure the fraction of this flux absorbed and re-radiated by dust. Each SFR indicator possesses its own strengths and disadvantages. The luminosity of the \\ha~recombination line is relatively simply and directly coupled to the incident number of Lyman continuum photons produced by young stars, and hence is proportional to the SFR \\citep{Kennicutt:1998pa}, although there is some dependence on the metallicity of the gas \\citep{Charlot:2001gj}. Provided the line can be adequately corrected for stellar absorption, the largest uncertainty is the correction for the presence of dust. At higher redshifts, $z\\gsim0.4$, the \\ha~line passes out of the optical window and the most commonly-used emission line in the range $0.4 \\lsim z \\lsim 1.5$ has been \\oii$\\lambda 3727$. \\oii~luminosity is more strongly affected by dust than \\ha. Moreover, this collisionally-excited line is very sensitive to metallicity. Since the emission lines necessary to accurately measure dust extinction and metallicity are not available when \\oii~is used (otherwise the \\ha~line would be used in preference to \\oii), alternate approaches to empirically correct the \\oii-SFR for these dependencies have been adopted \\citep[e.g.,][]{Moustakas:2006wp,Weiner:2007tf}. Another observational limitation of indicators based on spectroscopy, particularly for low redshift, is the need for an aperture correction. The spectroscopic aperture may only sample the innermost regions of the galaxy and thus an extrapolation is necessary to sample the total light from the galaxy. This may be further complicated by population gradients across the galaxy and the fixed angular size aperture sampling a different fraction of a given galaxy at different redshifts. A comprehensive examination of aperture effects within SDSS data is given in \\citet{Brinchmann:2004ct}. The rest-frame UV luminosity is straightforward to measure in high redshift surveys. Indeed, it is normally produced as a by-product of the survey data, since all that is required is a relatively blue optical band (depending on the redshift of the source). The UV luminosity provides a measurement of the continuum radiation from young stars, again after it has been extinguished by dust. Although sensitive to dust, the UV is relatively insensitive to metallicity (e.g., \\citealt{Glazebrook:1999kl}). Additionally, a correction must be made for the contribution to the UV luminosity from more evolved stellar populations. This correction becomes larger towards lower redshifts, as the Universe ages and so do the typical ages of the galaxies considered. The extinction due to dust may be estimated from the slope of the UV continuum (which is easily measured from two or more filter photometry). This method was originally applied, with considerable uncertainty, only to starburst galaxies \\citep{Meurer:1999qd}, but recent work has found comparable relations for normal galaxies (e.g., \\citealt{Cortese:2006rm}). The effect of dust on these indicators can be large and varies considerably from galaxy-to-galaxy. In the local Universe, the dust extinction in \\ha~can be 0-2 mag and in the UV 0-4 mag \\citep{Kennicutt:1998pa,Calzetti:2000sp}. This scatter may be partly driven by galaxy inclination and dust geometry effects, meaning that even for galaxies with the same dust content, a different effective attenuation may be observed. Furthermore, the different indicators are sensitive to the presence of young stars for different lengths of time: the ionising radiation responsible for the \\ha~line measures SFRs on timescales of $\\sim$10 Myr, whereas UV-bright stars which are considered in the SFR derived from UV luminosity have lifetimes $\\sim$100 Myr. Similarly, the sensitivity of the ultraviolet continuum and emission line luminosities to stars of different masses ($\\gsim5$ and $\\gsim10$ M$_\\odot$ respectively) means that the derived SFRs are sensitive to the IMF. This also means that if the assumption of a universal IMF is incorrect, this will lead to differences between these SFR indicators (e.g., \\citealt{Meurer:2009la}). Thus, the SFR is difficult to estimate accurately for an individual galaxy. One solution is to measure the SFRs for large ensembles of galaxies and compare their average properties in order to smooth over some of these effects. For example, comparing globally-averaged SFR quantities estimated from different indicators with different sensitivities to dust, metallicity, etc. should show the same results, on average, if the corrections for these effects are broadly consistent. Now, the potential systematic error of most concern is the timescale to which the different indicators are sensitive. However, provided the timescales of all the indicators used are short compared with the timescale over which the cosmic average SFR is changing, this difference can be negated. Since most of the variables affecting the estimate of SFR correlate with mass, e.g., metallicity via the mass--metallicity relation \\citep{Tremonti:2004ik}, dust \\citep{Brinchmann:2004ct}; an obvious first step is to determine mass-dependent corrections to the SFR indicators. Furthermore, there has been great activity recently in measuring how SFR (or some related quantity such as SFRD or SFR/stellar mass) varies as as a function of stellar mass, and how this evolves with redshift. Most of these results have been consistent with a picture in which the termination of star-formation progresses from higher mass to lower mass galaxies as the Universe ages, so-called {\\it cosmic downsizing} \\citep{Cowie:1996xw}. Obviously any residual systematic errors as a function of mass on the SFR indicators used could pose a serious problem when studying such relatively subtle mass-dependent effects. To date, most high-redshift studies have used locally-calibrated tracers of SFR, under the assumption that the calibration is still valid at higher redshift. Most local SFR calibrations use a simple conversion between luminosity of some SFR proxy and SFR and assume that this is valid for all masses. However, as we will show in this paper, a better approach is to allow for a mass-dependent conversion factor which accounts for differences such as the average dust extinction as a function of galactic mass. Dust obscuration at higher redshift is an open question. For example, in two recent z$\\sim$2 studies, \\citet{Pannella:2009uj} and \\citet{Reddy:2009wc} both find evidence for mass-dependent extinction, but the former find an extremely large dust correction is necessary (much larger than that measured locally, \\S~\\ref{sec:othemp}), whereas the latter find that the amount of dust extinction required is a {\\it decreasing} function of redshift. Part of this difference may be due to sample selection in that the \\citet{Reddy:2009wc} sample selects Lyman Break Galaxies, which may favour systems with lower UV extinction, but this still goes to show that the evolution of dust properties is not well constrained. Other works (e.g., \\citealt{Martin:2007xe,Iglesias-Paramo:2007db}) at z$\\sim$1 find similar evidence for mass-dependent extinction using the ratio of IR to FUV luminosity to estimate extinction. Regardless, a sensible null hypothesis would be to assume that the higher redshift relations (such as the mass dependence of extinction) follow the local relations, and derive SFR calibrations and corrections based on local data. Using a local mass-dependent correction at higher redshift is clearly preferable to ignoring the mass dependence of SFR indicators altogether. As a first step in understanding how the various SFR indicators commonly used at higher redshift relate to one another, we have undertaken a comparison between these indicators in the local Universe. The Sloan Digital Sky Survey (SDSS, \\citealt{york}) provides high quality optical photometry and spectra for a well-understood sample of galaxies. In addition, several groups have already performed detailed measurements of SFRs using state-of-the-art techniques (fitting spectrophotometric data to the latest stellar evolutionary synthesis and photoionisation models) which may be used as a reference against which simpler methods, such as those necessarily employed at higher redshift, may be compared. The main aim of this paper is to examine the consistency between different star-formation indicators by exploring the SFRD estimates obtained with the different tracers and to look for possible systematic differences as a function of mass. We examine these in the local Universe, but end with a brief exploration of how these may affect higher redshift surveys. This paper is designed to serve as a local reference, both for understanding how the various estimators compare with each other and to examine evolution when used in conjunction with higher redshift surveys. We proceed in the following way. In \\S2 we describe the SDSS optical imaging and spectroscopy and {\\it GALEX} UV photometry. \\S3 describes the SFR indicators examined. We use the \\ha~and \\oii~emission lines and UV photometry from the SDSS $u$-band and {\\it GALEX} FUV. For each of these indicators, we explore a range of different assumptions for the dust correction, e.g., a nominal constant extinction, extinction estimated from optical emission lines (or, for example, the UV slope). We present results for our `best estimator' for each indicator throughout the paper, but also show how these estimates change under the various different assumptions commonly made. Many of the prescriptions we use are well known in the literature (particularly the UV), but we will show that current \\oii~indicators are far from optimal and first we will need to correct this. In \\S4 we present our new empirical (mass-dependent) correction to \\oii; \\S5 examines the emission line luminosity functions and \\S6 presents the the SFRD results, convergence tests for the limiting SFR necessary to accurately estimate the SFRD, and also briefly considers the application of these to a higher redshift survey. A more detailed z$\\sim$1 analysis drawing on these local results will be presented in \\citet{Gilbank:2010hc}. In \\S7 we summarise our conclusions. All magnitudes are quoted on the AB system unless otherwise stated, and we assume a cosmology $(h,\\Omega_M, \\Omega_\\lambda)=(0.7,0.3,0.7)$. Throughout, we convert all quantities to those using a \\citet{Kroupa:2001ea} universal initial mass function (IMF), except for \\S\\ref{sec:highz}. ", "conclusions": "We have compared SFR indicators based on \\ha, \\oii, $u$-band and FUV luminosity using the SFRD as a function of stellar mass for $\\sim$50\\,000 in the SDSS Stripe 82 region ($0.032$2 are also the most dust obscured \\citep{Dole04,Papovich04,LeFloch09}. Nevertheless, mid- and far- infrared surveys (particularly those made with the 850$\\mu$m SCUBA camera on the JCMT and more recently with the {\\it Spitzer Space Telescope} at 24$\\mu$m) have begun to resolve the most highly obscured populations into their constituent galaxies, determine their contribution to the energy density in the extra-galactic far-infrared/sub-mm background, and chart that history of massive galaxy formation \\citep{Smail02,Cowie02,LeFloch05}. Extensive, multi-wavelength follow-up has shown that these heavily dust-obscured, gas-rich galaxies lie predominantly at high redshift ($z\\sim2$; e.g. \\citealt{Chapman03a,Chapman05a}), with bolometric luminosities of $\\gg$10$^{12}$L$_{\\odot}$ and star-formation rates of order 700\\,M$_{\\odot}$\\,yr$^{-1}$. It has therefore been speculated that SMGs are the progenitors of luminous elliptical galaxies (e.g.\\ \\citealt{Lilly99,Genzel03,Blain04a,Swinbank06b,Tacconi08}). With the redshift distributions and contributions to the cosmic energy density of ultra-luminous galaxies reasonably constrained, the next step is to study the evolutionary history of SMGs, and to determine how they relate to lower luminosity galaxies. Indeed, given the apparently rapid evolution in the space density of ULIRGs from $z\\sim2$ to $z=0$ \\citep{Chapman05a,LeFloch05}, one key issue is to understand the physical processes which trigger these far-infrared luminous events. Indeed, the mechanism responsible for these vigorous starbursts is still uncertain. Analogy to local ULIRGs would argue for merging as the trigger, although secular bursts in massive gas disks is also conceivable and indeed recent theoretical interest has stressed the importance of cold flows in high-redshift star formation \\citep[e.g.][]{Genel08}. The suggestion that SMGs have compact disk-like gas reservoirs ($R_{1/2}<2$\\,kpc) with ``maximal starbursts'' \\citep{Tacconi08} hints that SMGs are scaled-up versions of the local ultra-luminous galaxy population, which are usually associated with merger activity \\citep{Tacconi02}. In order to test the connection between SMGs, lower-luminosity star-forming galaxies at high redshift, as well as local ULIRGs, we have obtained high resolution {\\it Hubble Space Telescope (HST)} imaging of a sample of spectroscopically confirmed SMGs at $z=$0.7--3.4. By necessity, most morphological studies of high redshift galaxies to date have been performed at optical wavelengths which probe the rest-frame UV \\citep{Chapman03b,Webb03,Conselice03b,Conselice08,Smail04,Law07b}. However the rest-frame UV is dominated by radiation which traces the brightest, active star-forming regions rather than the bulk of the stellar population and can lead to late-type galaxies being classified as irregular systems \\citep{Dickinson00,Thompson03,Goldader02}. Nevertheless, in a recent study, \\citet{Law07b} conducted a detailed analysis of the rest-frame UV morphologies of a large sample of UV/optically selected star-forming galaxies at $z\\sim$1--3 in the Hubble Deep Field North (HDFN) and find evidence that dusty galaxies have more nebulous UV morphologies than more typical sources, but otherwise conclude that UV morphology is statistically decoupled from the majority of physical observables (such as stellar or dynamical mass, gas fraction or star-formation rate). Here we aim to extend this work to include the rest-frame optical emission to test whether there are key differences in the morphologies at longer wavelengths (as suggested by imaging of low redshift ULIRGs; e.g. \\citealt{Goldader02}). We have therefore assembled a sample of 25 SMGs with both {\\it HST} ACS $I$-band and NICMOS $H$-band observations. We determine the basic morphological parameters of this sample of SMGs, as well as search for signs of tidal features and major mergers. To baseline our analysis we also use a sample of 228 optically selected star-forming galaxies at $z\\sim2$--3 (of which 53 have also been observed in $H$-band with {\\it HST}). We use a {\\it WMAP} cosmology \\citep{Spergel04} with $\\Omega_{\\Lambda}$=0.73, $\\Omega_{m}$=0.27, and H$_{0}$=72\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. In this cosmology, at $z=$2.1 (the median redshift of our SMG sample), 0.1$\"$ (the typical resolution of our observations) corresponds to a physical scale of 0.8\\,kpc. All quoted magnitudes are on the AB system unless otherwise noted. \\begin{figure*} \\centerline{\\psfig{file=figs/col.ps,angle=0,width=6.0in}} \\caption{High resolution {\\it HST} optical (ACS/WFPC2/STIS) and near-infrared NICMOS imaging of SMGs. For each galaxy, the left hand panel denotes the optical image whilst the right hand panel denotes the 1.6$\\mu$m ($H-$band) image. Each thumbnail is scaled such that they each cover 40 kpc at the redshift of the SMG. The small cross denotes the location of the centroid of the radio emission.} \\label{fig:hstthumbs1} \\end{figure*} \\begin{figure*} \\centerline{\\psfig{file=figs/colSMG.ps,angle=0,width=6.0in}} \\caption{True colour {\\it HST} $IH$-band images of the SMGs in our sample showing the range of colours and morphological mix within the sample. Each image is 40\\,kpc at the redshift of the galaxy.} \\label{fig:colSMG} \\end{figure*} \\begin{figure*} \\centerline{\\psfig{file=figs/col_LBG.ps,angle=90,width=6.4in}} \\caption{Comparison true colour {\\it HST} $IH$-band images of thirty star-forming galaxies at $z\\sim2$ from the spectroscopic sample of \\citet{Reddy06b} showing the mix of morphological types is comparable to the SMGs. As in Fig.~\\ref{fig:hstthumbs1}, each image is 40\\,kpc at the redshift of the galaxy so that a direct comparison can me made.} \\label{fig:colLBG} \\end{figure*} \\begin{figure*} \\centerline{ \\psfig{file=figs/rh_z.ps,width=6in,angle=90}} \\caption{Size versus redshift relation for SMGs (filled red circles) compared to the UV-SF (filled blue square) and high-$z$ field (small black squares) comparison samples. The half light radii in the left hand panel are derived using {\\it HST} $I$-band data whilst the right hand panel are derived from the NICMOS $H$-band (F160W) imaging. The dashed lines show that the medians of the distribution and illustrate that the SMGs and comparison samples have comparable half light radii in both the $I$- and $H$-bands. On the right hand side of each sub-panel we show the cumulative histograms of each distribution which shows that in the rest-frame UV the SMGs have marginally larger half light radii on average than the comparison samples, but essentially indistinguishable distributions in the rest-frame optical.} \\label{fig:rh_z} \\end{figure*} \\begin{figure} \\centerline{ \\psfig{file=figs/Gini_Asym_opt.ps,angle=90,width=3.0in}} \\caption{Comparison between the structural parameters for SMGs compared to star-forming field galaxies. {\\it Top:} Gini versus Asymmetry measured in the observed $I$-band for the SMGs showing that the median asymmetry for the SMGs and comparison samples is comparable, but that there is an offset between the Gini co-efficients of $\\Delta G\\sim$0.1 between SMGs and the comparison samples. The dashed line illustrates the correlation between Gini and Asymmetry in local galaxies (E-Sd morphological types) from Lotz et al. (2004). The cumulative histograms on each axis show that there is a significant difference in the rest-frame UV Gini coefficient for the SMGs compared to the comparison samples, but the asymmetries are comparable to the comparison samples. {\\it Centre:} $H$-band Gini versus Asymmetry for the SMGs and field star-forming galaxies. As in the observed $I$-band the asymmetry is indistinguishable between the populations. However, the median Gini coefficient for the SMGs in only $\\sim$0.03 larger than the field population. The dashed line denotes the same correlation as shown in the top panel. The cumulative histograms show that there is a subtle difference between the Gini co-efficient in the rest-frame optical for the SMGs and comparison samples, but the asymmetries are indistinguishable. {\\it Bottom:} $\\Delta G$ versus $\\Delta A$ for the SMGs and field galaxies showing that the both the SMGs and field populations tend to prefer larger $\\Delta G$ values.} \\label{fig:GiniAsym} \\end{figure} ", "conclusions": "We have undertaken the first large near-infrared morphological analysis of SMGs. % We find that the SMGs have comparable sizes to UV-SF galaxies at the same epoch (BX/BM and LBGs) in both $I$- and $H$-bands, and (surprisingly) with comparable asymmetries. However, we find that the SMGs have systematically larger Gini co-efficients (particularly in the observed $I$-band) than UV-SF galaxies at the same epoch suggesting less uniform, high intensity star-formation in the rest-frame UV, possibly reflecting structured dust obscuration (see also \\citealt{Law07b}). % Overall our results suggest that SMGs are no more likely to appear as major mergers in the rest-frame UV/optical than more typical, lower-luminosity high-redshift galaxies (such as Lyman break galaxies or BX/BMs). However, the differences in the Gini coefficients between populations suggests that the dustier SMGs have star-formation which is less uniform (and/or that they suffer from more structured dust obscuration). We also show that most of the SMGs have observed $H$-band light profiles which are better fit with that of a spheroid galaxy light distribution. Indeed, stacking the galaxies we find that in the observed $H$-band, an n$\\sim$2 Sersic index provides a better fit to the spatial light profile than an exponential disk model, suggesting that, whilst these galaxies are individually morphologically complex, the composite stellar structure of the SMGs reflects that of a spheroid/elliptical galaxy. However, we note that the same analysis of the UV-SF galaxies is statistically indistinguishable, with only a marginally lower Sersic index with $n=1.6\\pm0.3$. The close similarity between the rest-frame UV and optical morphologies in the SMGs, and UV-SF galaxies suggests that both wavelengths are dominated by young, star-bursting components as well as dusty regions \\citep[see also][]{Dickinson00,Papovich05}. These results are in contrast to local studies of similarly luminous LIRGS and ULIRGs in the local Universe which have been shown to have very different Hubble types from the rest-frame UV and optical wavelengths (e.g. \\citealt{Goldader02}), possibly suggesting fundamental differences between starbursts at $z=0$ and $z\\sim2$. Although the sizes and structural properties of the SMGs and UV-SF star-forming galaxies are similar, previous work has shown that the dynamical masses of SMGs are up to an order of magnitude larger than UV-SF galaxies \\citep[e.g.][]{Erb03,Swinbank04,Erb06a,Genzel06,ForsterSchreiber06,Swinbank06b,Law07}.. This suggests that the intense star-formation within SMGs does not represent an evolutionary sequence in which typical UV-SF galaxies undergo intense star-formation (either through merging or through secular processes), but rather it is the availability of large gas reservoirs within already massive galaxies that allow the SMG phase. Finally, we also investigate the size--stellar mass relation of SMGs in order to test whether the SMGs may represent progenitors of the luminous red galaxies seen at $z\\sim$1.5. We combine estimates of the size and stellar masses to show that approximately half of the sample have stellar mass densities comparable to those derived for luminous red galaxies \\citep[eg.][]{Cimatti08}. However, we show that the median size of the SMGs in the observed near-infrared ($r_{h}=2.3\\pm0.3$\\,kpc) is larger than that of the luminous red galaxies, which have a median half light radius of $r_{h}=1.2\\pm0.2$\\,kpc. We also show that the median stellar mass of the SMGs is also a factor $\\sim2\\times$ larger, thus suggesting that the luminous red population at $z\\sim1.5$ are unlikely to be direct descendants of the SMG population unless the new stars formed in the SMG starburst ultimately have a very different spatial distribution from their gas reservoirs, which seems unlikely given the similarity between the CO and UV/optical sizes. Overall, our results suggest that rest-frame UV and optical morphologies of high-redshift galaxies are essentially decoupled from other observables (such as bolometric luminosity, stellar or dynamical mass). Alternatively, the physical processes occuring within the galaxies are too complex to be simply characterised by the rest-frame UV/optical morphologies. It may be that at significantly longer wavelength structural differences will appear, but this will have to wait for rest-frame near-infrared imaging with {\\it James Webb Space Telescope, (JWST)}. Alternatively, high resolution kinematical studies on sub-kpc scales (e.g. with near-infrared IFUs which can now be carried out from the ground using adaptive optics) may offer the most direct route to probing the differences between high-redshift galaxy population as the dynamics, distribution of star-formation and metallicity gradients will reflect differences in the triggering mechanism and mode of star-formation at high-redshift." }, "1002/1002.0806_arXiv.txt": { "abstract": "As part of our comprehensive long-term multi-waveband monitoring of 34 blazars, we followed the activity in the jet of the blazar PKS~1510$-$089 during major outbursts during the first half of 2009. The most revealing event was a two-month long outburst that featured a number of $\\gamma$-ray flares. During the outburst, the position angle of optical linear polarization rotated by about $720^\\circ$, which implies that a single emission feature was responsible for all of the flares during the outburst. At the end of the rotation, a new superluminal knot ($\\sim 22c$) passed through the ``core'' seen on 43 GHz VLBA images at essentially the same time as an extremely sharp, high-amplitude $\\gamma$-ray and optical flare occurred. We associate the entire multi-flare outburst with this knot. The ratio of $\\gamma$-ray to synchrotron integrated flux indicates that some of the $\\gamma$-ray flares resulted from inverse Compton scattering of seed photons outside the ultra-fast spine of the jet. Because many of the flares occurred over time scales of days or even hours, there must be a number of sources of IR-optical-UV seed photons --- probably synchrotron emission --- surrounding the spine, perhaps in a slower sheath of the jet. ", "introduction": "There are four primary methods for probing the structure and physics of relativistic jets in blazars on parsec and sub-parsec scales: VLBI imaging in both total and polarized intensity at millimeter wavelengths, variability of the flux at radio through $\\gamma$-ray frequencies, polarization at radio through optical wavebands, and the spectral energy distribution (SED). In the past, efforts to study this capricious class of objects have been limited by datasets with substantial gaps in either time or frequency coverage. Starting in the mid-1990s, this situation has been greatly alleviated by the availability of instruments such as the Very Long Baseline Array (VLBA), the {\\it Rossi} X-ray Timing Explorer (RXTE), and, most recently, the {\\it Fermi} Gamma-ray Space Telescope. Together with concerted efforts with optical/near-infrared, millimeter-wave, and radio telescopes, these facilities now allow long-term, comprehensive multi-waveband monitoring of a number of blazars. Such programs are providing valuable insights into the processes by which jets form and propagate, as well as the locations and physics of flux outbursts in blazars \\citep[e.g.,][]{Mar08,Chat08,Lar08}. We are leading a monitoring program that provides the data needed to combine the techniques listed above in order to locate the sites of high-energy emission, determine the processes by which X-rays and $\\gamma$-rays are produced, and infer the physical conditions of the jet, including the geometry of its magnetic field. Because of limitations in sensitivity and available observing time, such comprehensive monitoring can be mounted only for a relatively small number of objects, 34 in the case of our program. Progress toward an overall understanding of blazars therefore involves generalization of inferences drawn from the well-studied specimens into a framework that can be applied to the entire class. We present here a rich set of observations of the quasar PKS~1510$-$089 ($z$ = 0.361) that allows us to locate the sites of $\\gamma$-ray flares, leading to inferences on the high-energy emission processes. PKS~1510$-$089 possesses all of the characteristics of blazars: flat radio spectrum, apparent superluminal motion --- among the fastest \\citep[as high as $45c$ for H$_\\circ$=71 km s$^{-1}$Mpc$^{-1}$ and the concordance cosmology][]{Spergel07} of all blazars observed thus far \\citep{J05} --- high-amplitude and rapid flux variability at all wavebands, strong and variable optical linear polarization, and high $\\gamma$-ray apparent luminosity \\citep{Hartman99}. The observations of this blazar from our program allow us to compare motions, linear polarization, and changing flux of features in the parsec-scale radio jet with flux variability of the entire source at radio, millimeter-wave, IR, optical, X-ray, and $\\gamma$-ray frequencies, and with optical polarization. The relative timing of the correlated variations thus found probe the structure and physics of the innermost jet regions where the flow is accelerated and collimated, and where the emitting electrons are energized. ", "conclusions": "We are finally at the point where the richness of our datasets is sufficient for us to draw grand inferences about the locations and physical mechanisms of $\\gamma$-ray flares. If PKS~1510$-$089 and BL~Lac are typical, it is no wonder that previous, less comprehensive monitoring programs left us confused! There are multiple sources of seed photons, some of which are quite local (as opposed to more global, such as an IR-emitting hot dust torus) and may not radiate at sufficiently high luminosities to be directly observable. These include the accretion disk and BELR --- which may be important during the earliest part of a multi-flare outburst --- shocks or other features in the sheath of the jet, and synchrotron radiation from the fast spine of the jet that contains the electrons that scatter the seed photons to $\\gamma$-ray energies. A single superluminal knot can be responsible for a number of flares and periods of sustained elevated $\\gamma$-ray emission. As we accumulate data for more blazars, we expect to see similar patterns of flares both before and after a disturbance passes through the core of the jet. We also hope to find deviations from this pattern that serve to provide further insight into the range of physical behavior of blazars. In both cases, we anticipate a tremendous increase in our understanding of the relativistic jets of blazars by combining the great power of monitoring the flux, polarization, and sub-milliarcsecond structure of these exciting objects. \\vspace{-5mm}" }, "1002/1002.5048_arXiv.txt": { "abstract": "{The Rayleigh-Taylor instabilities that are generated by the deceleration of a supernova remnant during the ejecta-dominated phase are known to produce finger-like structures in the matter distribution that modify the geometry of the remnant. The morphology of supernova remnants is also expected to be modified when efficient particle acceleration occurs at their shocks.} {The impact of the Rayleigh-Taylor instabilities from the ejecta-dominated to the Sedov-Taylor phase is investigated over one octant of the supernova remnant. We also study the effect of efficient particle acceleration at the forward shock on the growth of the Rayleigh-Taylor instabilities.} {We modified the Adaptive Mesh Refinement code RAMSES to study with hydrodynamic numerical simulations the evolution of supernova remnants in the framework of an expanding reference frame. The adiabatic index of a relativistic gas between the forward shock and the contact discontinuity mimics the presence of accelerated particles.} {The great advantage of the super-comoving coordinate system adopted here is that it minimizes numerical diffusion at the contact discontinuity, since it is stationary with respect to the grid. We propose an accurate expression for the growth of the Rayleigh-Taylor structures that smoothly connects the early growth to the asymptotic self-similar behaviour.} {The development of the Rayleigh-Taylor structures is affected, although not drastically, if the blast wave is dominated by cosmic rays. The amount of ejecta that reaches the shocked interstellar medium is smaller in this case. If acceleration were to occur at both shocks, the extent of the Rayleigh-Taylor structures would be similar but the reverse shock would be strongly perturbed.} ", "introduction": "\\label{int} In young supernova remnants (hereafter SNR), the dense shell of material ejected by the explosion and decelerating in a rare\\-fied interstellar medium (hereafter ISM) is expected to be subject to hydrodynamic instabilities \\citep{g73,g75,s78} of Rayleigh-Taylor (hereafter RT) type. These instabilities modify the morphology of the SNR causing a departure of the ejecta from spherical symmetry. They manifest themselves as finger-like structures of material protruding from the contact discontinuity between the two media into the ISM heated by the forward shock, as shown by numerical simu\\-lations \\citep[e.g.][]{cbe92,dc98,d00,wc01}. During this process, the shocked ejecta and the shocked ISM remain two distinct fluids. The X-ray observations of Tycho's SNR \\citep{w05} exhibit structures that are consistent with these effects. Deviations from spherical symmetry may also be caused by initial asymmetries in the explosion of the progenitor or local inhomogeneities in the circumstellar medium. In Cassiopeia A, the spatial inversion observed by the Chandra observatory of the iron and silicon layers \\citep{hrbs00} provides a strong indication of an asymmetric explosion. An even more radical example of deviation from spherical symmetry is SN 1993J, a stellar wind case, where the optical and radio observations can be reconciled by assuming a strong asphericity for the pre-existing progenitor activity \\citep{bbr01}. Two-dimensional hydrodynamic simulations of SNRs can take into account the RT instability. Three-dimensional simulations have shown an enhancement of small-scale structures and more severe deformation of the reverse shock surface \\citep{jn96}; the perturbation seems to grow faster by $30\\%$ than in the two-dimensional case \\citep{k00}. We chose to pursue three-dimensional hydrodynamic numerical simulations, focusing on the deviations from spherical symmetry of the SNR in the ejecta-dominated phase induced by the RT instabilities. Several distinct physical processes occur in young SNRs, for instance, the dispersion of synthesized materials in the circumstellar medium and the propagation of collisionless shocks. Galactic SNRs are also considered to be a strong candidate source of galactic cosmic rays up to the knee, i.e. $E \\sim 3\\times10^{15}$ eV \\citep{lc83}, since the rate and energy budget can account for the galactic energy density of cosmic rays; however, a direct identification of SNRs as sources of cosmic-ray nuclei is still lacking. In the present paper, we aim to investigate with hydrodynamic equations the growth of Rayleigh-Taylor instabilities in the context of efficient particle acceleration. We adapt to SNR physics the hydrodynamic version of the AMR (Adaptive Mesh Refinement) code RAMSES \\citep{t02}, designed originally to study the large-scale structure formation in the universe with high spatial resolution. The first application, considered here, is to study the effect of the growth of RT instabilities on the profile of the hydrodynamic variables by considering the evolution of a full octant of the SNR, i.e. a larger angular region than previously considered \\citep{be01}. We do not introduce any seed perturbation into the 3D radial velocity field, so that any departure from spherical symmetry is naturally produced by the numerical fluctuations; this is supported by the non-linear phase of the instability being insensitive to initial perturbations \\citep{cbe92}. \\citet{be01} described the efficient acceleration of particles by changing the effective adiabatic index throughout. A higher compression ratio at the shocks was found to only slightly affect the growth of RT instabilities. However, the suggestion that the reverse shock can efficiently accelerate particles continues to be debated \\citep{edb05}. In this paper, we study the impact of cosmic-ray acceleration on the instabilities using a simplified description that assumes that the adiabatic index is 4/3 in the shocked ISM region, but remains 5/3 inside the contact discontinuity (surface separating ejecta and ISM). This simulates the case of a cosmic ray-dominated blast wave and gas-dominated reverse shock. To follow the expansion of SNRs, we use synergically two numerical approaches: AMR and the Moving Computational Grid. As a result, the large-scale turbulence in SNR can be simu\\-lated, allowing for an accurate description of the instabilities. The hydrodynamic equations can be formulated with respect to two distinct classes of coordinate systems: Eulerian, i.e., fixed space-coordinates in time, whose major concern is the production of numerical diffusion due to the advective terms; and Lagrangian, i.e., coordinates comoving with the bulk fluid, which is free in principle from numerical diffusion, but possible grid distortions require a rezoning which produces new numerical diffusion. Therefore, Eulerian coordinate systems are generally selected for multidimensional flow simulations. The numerical approach of Moving Computational Grid consists of using a computational grid comoving, or quasi-comoving, with the hydrodynamic flow to minimize the local fluid velocity. The idea of a computational grid adapted to follow the global motion of the fluid was first proposed in cosmological numerical simu\\-lations by \\citet{g95}. However, further analysis \\citep{gb96} highlighted problems in the coupling of the hydrodynamics solver with the gravity solver. Strong mesh deformations were also found in the gravitational clustering. In the approach of Moving Computational Grid, the mesh moves continuously and in addition to a full transformation of coordinates (position and time), a transformation of hydrodynamic variables (density, velocity and pressure) is performed. In contrast, in a purely Lagrangian approach the space-coordinates are comoving with the bulk flow, leaving the hydrodynamic quantities unchanged. The main disadvantage of simulating a SNR in a Eulerian fixed computational grid is that in the bulk flow of the SNR the total energy is roughly equal to the kinetic energy. Therefore, the thermal energy, computed as the difference between total and kinetic energies, can be very small leading possibly to local negative thermodynamic pressure. An algorithm is required to control the sign of pressure. We apply a combination of the AMR with the Moving Computational Grid from the {\\it young} phase, starting soon after the self-similar profile of \\citet{c82,c83} has been established. The modification introduced enabled us to follow the evolution of one eighth of the volume of a young SNR, a large volu\\-me with respect to previous 3D simulations \\citep[e.g.][]{jn96}, and for a long interval of time, namely until the transition to the Sedov-Taylor phase. We do not provide all the details of the RAMSES code, which can be found in \\citet{t02}, but a description of the modifications introduced here is outlined. The plan of the paper is the following. In Sect.~\\ref{inicond}, we discuss the initial conditions adopted for a young SNR, namely the mapping of the uni-dimensional solution of the hydrodynamic equation over a cartesian grid. In Sect.~\\ref{eqs}, we present the hydrodynamic equations for the SNR flow, written in both the laboratory frame and the reference frame which is comoving with the contact discontinuity. In Sect.~\\ref{num}, we sketch the main steps of the implementation of the Godunov method in RAMSES, we discuss the modifications and the new variables introduced: the active variable $\\alpha$ to change locally the equation of state; the passive scalar $f$, which traces the surface of the contact discontinuity; and the ionization age $\\tau$. In Sect.~\\ref{snr}, we present the results of SNR simulations: the growth of the RT structures and the elongation in time; the influence of a cosmic-ray-dominated blast wave on the RT instabilities. In Sect.~\\ref{disc}, we discuss the main findings and present additional numerical tests. In Sect.~\\ref{conc}, we conclude and present forthcoming applications. ", "conclusions": "\\label{conc} We have adapted the AMR code RAMSES, which is based on a second order Godunov method in an expanding frame called ``supercomoving coordinates'', to follow the evolution of SNRs. In this approach, not only the space-time variables have been modified but also the hydrodynamics variables. The comoving coordinate system allows an eighth of the total volume of the SNR to be described, which is much larger than previously considered. A longer time interval can be investigated, of the order of thousands of years, until the transition to the asymptotic Sedov-Taylor phase. Such a large volume allows the convective instabilities to be modeled more accurately since it considers not only the shortest wavelengths, which have the greatest growth rate, but also the longer wavelengths, which grow more slowly but make a significant contribution to the morphology of the SNR over time-scales of thousands of years. Our larger spatial sampling allows a statistically more accurate description of the instabilities. The great advantage of the method adopted here is that it minimizes the velocity of the fluid relative to the grid, and the numerical diffusion at the contact discontinuity, where the instability should develop. In the comoving reference frame, the contact discontinuity, in the absence of distortions due to convective instabilities and before the transition to the Sedov-Taylor phase, would be exactly stationary in time. The analytical non-inertial term resulting from this coordinate transformation is strictly equivalent to a gravitational acceleration: the Euler equations have therefore been integrated with the corresponding additional source terms. The elongation of the Rayleigh-Taylor structures slowly reaches the asymptotic behaviour $t^\\lambda$, by direct solution of the self-similar theory according to the assumption of a self-similar acceleration instead of a constant acceleration. We have presented a simple way to numerically investigate the effect of efficient particle acceleration at the forward shock, approximating the shocked ISM as a relativistic gas and the shocked ejecta as a non-relativistic gas, with a different adiabatic exponent in each fluid. The density behind the forward shock may be higher than at the reverse shock in that case. A deceleration in the protruding of RT structures is caused by the higher compression of the shocked ISM. The conclusion of \\citet{be01} that the RT fingers can travel very close to the shock in the presence of accelerated particles remains valid. The elongation of the instabilities has here been more precisely quantified. We propose that this will allow us to understand why ejecta are found very close to the blast wave in the remnant of SN 1006 \\citep{ch08}. The set-up of the code presented here will allow future studies of the back-reaction of particle acceleration on the SNR evolution." }, "1002/1002.2160_arXiv.txt": { "abstract": "Star clusters are fundamental building blocks of galaxies. Their formation is related to the density and pressure in progenitor molecular clouds and their environmental conditions. To understand better the dynamical processes driving star formation and chemical evolution, we compare ages, metallicities, and alpha-element abundance ratios of globular clusters in nearby dwarf galaxies of different luminosities and morphological types, and situated in different environments. The data are based on our 6m telescope medium-resolution spectroscopic observations. We find that a mean metallicity of GCs in a galaxy at a given age is higher for early-type dwarfs, than for late-type dwarf irregulars and spirals. ", "introduction": "Low surface brightness (LSB) dwarf galaxies are fainter than $M_B=-16^m$ and they represent the most common type of objects in the local Universe. In distinction to globular clusters (GCs), they are dark-matter dominated. Dwarf spheroidals (dSph) are gas-poor, old stellar systems. Dwarf irregulars (dIrr) contain young stars and neutral hydrogen, the fuel for star formation activity. Some objects, classified as transitional types (dIrr/dSph), show evidence for a few young and intermediate-age stars but not for any gas, or contain very small amounts of it, which is below the detection level of most modern telescopes. The origin of dwarf galaxies is under debate. What processes drive morphological transformations? Why do some of these objects lose their gas? Clues to the solution of these problems can be found in the observational properties of these objects and their GC systems. Ancient GCs are the brightest representatives of the oldest simple stellar populations. Their younger cousins, massive compact star clusters, have formed in some galaxies over the lifetime of the universe. The physical properties of both (hereafter: GCs) are often used for a better understanding of the evolution of dwarf galaxies and their role as building blocks in the cosmological structure formation picture (see for a review West et al., 2004, and references therein). \\medskip ", "conclusions": "We discuss the PDFs of evolutionary parameters (age, metallicity, [$\\alpha$/Fe]) of GCs observed in our 6m telescope spectroscopic campaign and compare those with GCs from the literature. Our analysis shows that 1) the metallicity and age dispersions in GC systems are wider for larger, i.e. more massive galaxies; 2) metal-rich clusters are preferentially found in galaxies more massive than $\\sim\\!10^9 M_{\\sun}$; 3) intermediate-age GCs in early-type dwarf galaxies are richer in metals than star clusters representing dynamically \"cold\" gas-rich environments in dIrrs; \\begin{figure}[!h] \\centerline{\\hbox{\\psfig{figure=sharinaFig1.ps,angle=0,clip=,width=13cm}}} \\caption[]{Probability density functions of evolutionary parameters (age, metallicity, [$\\alpha$/Fe]) for GCs in our Galaxy, M31, LMC, and dwarf galaxies in nearby groups. The curves indicate non-parametric probability density estimates using an Epanechnikov kernel (see text for details). } \\label{fig1} \\end{figure}" }, "1002/1002.2951_arXiv.txt": { "abstract": "Between 2004 and 2009 a sample of 28 X-ray selected high- and intermediate-frequency peaked blazars with a X-ray flux larger than 2\\,$\\mu\\rm{Jy}$ at 1\\,keV in the redshift range from 0.018 to 0.361 was observed with the MAGIC telescope at energies above 100\\,GeV. Seven among them were detected and the results of these observations are discussed elsewhere. Here we concentrate on the remaining 21 blazars which were not detected during this observation campaign and present the 3\\,sigma (99.7\\,\\%) confidence upper limits on their flux. The individual flux upper limits lie between 1.6\\,\\% and 13.6\\,\\% of the integral flux from the Crab Nebula. Applying a stacking method to the sample of non-detections with a total of 394.1 hours exposure time, we find evidence for an excess with a cumulative significance of 4.9 standard deviations. It is not dominated by individual objects or flares, but increases linearly with the observation time as for a constant source with an integral flux level of $\\sim$1.5\\,\\% of that observed from the Crab Nebula above 150\\,GeV. ", "introduction": "MAGIC (\\textit{M}ajor \\textit{A}tmospheric \\textit{G}amma-ray \\textit{I}maging \\textit{C}herenkov) is currently a system of two 17\\,m telescopes located atop the Roque de los Muchachos on the Canary Island of La Palma at 2200\\,m a.\\,s.\\,l. The observations refered to in this study were obtained during the years 2004 - 2009 when MAGIC was still a single-dish telescope. Its 234\\,$\\rm{m}^2$ tessellated parabolic mirror allows observations of VHE (\\textit{V}ery \\textit{H}igh \\textit{E}nergy) $\\gamma$-rays between $\\sim$50\\,GeV and 10\\,TeV. One key goal of the MAGIC telescope project is to determine the properties of extragalactic VHE sources, among which the high-frequency peaked BL Lacertae objects are the most numerous. Blazars are a subclass of radio-loud AGN and belong to the most extreme and powerful objects in the universe. They are characterized by a non-thermal broad-band continuum emission which is highly variable on time scales from years down to minutes \\citep{alb2007,aha2007a}. The spectral energy distribution (SED) of blazars is characterized by two bumps in a $\\nu$\\,F$_{\\nu}$ representation. The first component peaks at energies between IR and hard X-rays, and is assumed to originate from leptonic synchrotron radiation. The maximum of the second peak lies in the $\\gamma$-ray energy regime. The origin of this peak can be explained by different and partially concurring models either relying on inverse Compton scattering of electrons \\citep{mar,der,sik} or proposing hadronic interactions inside the jet \\citep{man,pro}. In the case the synchrotron peak occurs at energies above $\\sim10^{16.5}$\\,Hz, \\citep[according to][]{nie} these blazars are called HBLs (high-frequency peaked BL Lacertae objects) and for peak energies of $\\sim10^{14.5-16.5}$\\,Hz IBLs (intermediate BL Lacertae objects). As of April 2010, altogether 29 blazars were established as VHE sources (24 of them HBLs including M87 as 'misaligned' blazar)\\footnote{cf.\\ http://wwwmagic.mppmu.mpg.de/$\\sim$rwagner/sources/ for an up-to-date list.}, compared to six HBLs, when the MAGIC telescope began its regular observations in December 2004. The sample presented here comprises 21 X-ray selected objects which were not detected in the VHE regime prior to the MAGIC observations. Nine of the objects were already observed between December 2004 and February 2006 and the upper limits of these observations are reported in \\citet{alb2008a}. As there have been improvements within the MAGIC analysis, the data of these objects were re-analyzed and the new results are presented in this work. Since no significant detection was attained, upper limits on a 3\\,$\\sigma$ (99.7\\,\\%) confidence level will be presented. None of the observed sources showed any variability on diurnal timescales in the VHE regime. Assuming a positive detection in the case of a flaring state, the observations presented here provide a means of investigating the baseline emission of these objects. Therefore, a stacking method applied to the blazar sample can reveal such an emission below the sensitivity limit for each individual object. Together with VERITAS \\citep{ben2009} this is the second stacking analysis which turns out to be successful in the VHE $\\gamma$-ray regime. Former experiments like HEGRA failed in detecting a significant signal in a stacking analysis due to their limited sensitivity \\citep[cf.\\ for instance][]{man96}. In Section \\ref{sec:sample} the selection criteria for the objects will be presented. The observations and the data analysis technique are described in Section \\ref{sec:obsanal}. The analysis results are shown in Section \\ref{sec:results}. Finally, a discussion of the results and inherent implications can be found in Section \\ref{sec:disc}. ", "conclusions": "In the course of the MAGIC observational program during 2004 - 2009, a major part was spent on X-ray bright BL Lacertae objects. For 21 non-detections upper limits on the integral flux ranging between 1.6\\,\\% and 13.6\\,\\% of the Crab Nebula flux could be determined. Applying a stacking method to the individual non-detections we found an average VHE emission of the sample of X-ray selected blazars at the 4.9\\,$\\sigma$ significance level above 100\\,GeV. It turns out out that the mean VHE $\\gamma$-ray flux is significantly lower than in archival X-ray measurements. The two-point spectral index between 1\\,keV and 200\\,GeV is 1.09$\\pm$0.04." }, "1002/1002.1213_arXiv.txt": { "abstract": "\\citeauthor{lk73} showed that parallax measurements are systematically overestimated because they do not properly account for the larger volume of space that is sampled at smaller parallax values. We apply their analysis to neutron stars, incorporating the bias introduced by the intrinsic radio luminosity function and a realistic Galactic population model for neutron stars. We estimate the bias for all published neutron star parallax measurements and find that measurements with less than $\\sim 95\\%$ certainty, are likely to be significantly biased. Through inspection of historic parallax measurements, we confirm the described effects in optical and radio measurements, as well as in distance estimates based on interstellar dispersion measures. The potential impact on future tests of relativistic gravity through pulsar timing and on X-ray--based estimates of neutron star radii is briefly discussed. ", "introduction": "\\label{sec:intro} The past decade has seen an exponential increase in the number of parallax measurements to radio pulsars. Whereas only 14 such measurements were known by the end of the year 2000, currently 52 radio pulsars have their distance determined either through VLBI or pulsar timing measurements (see Tables \\ref{tab:curpxs} and \\ref{tab:oldpxs}). The importance of radio pulsar parallaxes arises from the wide variety of highly accurate investigations to which pulsars lend themselves. For example, the combination of an accurate parallax with the pulsar dispersion measure (DM: the integrated electron density between the pulsar and Earth) provides an average electron density measurement along the line of sight, which can be used to construct Galactic electron density models \\citep[see, e.g.,][]{cl02}. Also, the highly polarised nature of pulsars allows measurement of the Faraday rotation which - in combination with a distance - can be used to map out the component of the Galactic magnetic field parallel to the line of sight \\citep{hml+06}. Finally, in pulsar timing, distances are essential in correcting spindown and orbital period derivatives for the Shklovskii effect caused by proper motion \\citep{shk70}. Consequentially, certain pulsar timing tests of general relativity are dependent on accurate distance measurements, as described by \\citet{dt91} and used in \\citet{nss+05} and \\citet{dvtb08}, amongst others. The large increase in the number of radio pulsar parallaxes and their unique applications warrant an investigation into potential biases. As pointed out by \\citet{lk73}, in an homogeneous field of stars the number of stars per unit of distance increases as $D^2$, where $D$ is distance. This makes it statistically more likely to find an object at larger distances where more volume is sampled and, hence, more sources lie. \\citet{lk73} showed analytically that this statistical underestimate of the stellar distance depends on the precision of the parallax measurement. Specifically, they derived the following proportionality: \\begin{equation} p(\\varpi | \\varpi_{\\rm 0}) \\propto \\left( \\frac{\\varpi_{\\rm 0}}{\\varpi}\\right)^4 \\exp \\left( - \\frac{\\left(\\varpi-\\varpi_{\\rm 0}\\right)^2}{2 \\sigma^2}\\right), \\end{equation} where $\\varpi_{\\rm 0}$ is the measured parallax, $\\sigma$ is the standard deviation of the measurement and $p(\\varpi | \\varpi_{\\rm 0})$ is the probability distribution of the actual parallax, $\\varpi$, given the measurement. \\citet{bm98} expanded this analysis by using the intrinsic luminosity function for the star's spectral class as a further source of prior information to use in the estimate of a more accurate parallax value. Two points of confusion have pervaded the literature on the topic of Lutz-Kelker bias. First, the effect we describe as Lutz-Kelker bias is commonly referred to as Malmquist bias in extragalactic astronomy. As \\citet{gf97} point out, the Malmquist bias originally referred to a positive bias in luminosity of magnitude-limited samples and has only recently acquired the meaning of the geometric bias described by \\citet{lk73}. \\citet{smi03b} stresses a second point of confusion and difference between the original Malmquist bias and Lutz-Kelker bias, namely the fact that the bias described by \\citet{lk73} and discussed in the present paper, refer to individual parallax measurements rather than to the overall average of a sample of objects. In this paper we revisit the analysis by \\citet{bm98} with a particular focus on radio pulsars. To account for the analytically complex but realistic pulsar luminosity function and Galactic pulsar distribution, we describe a Monte-Carlo approach to correct previously published parallax measurements for the Lutz-Kelker bias. Our basic analytic derivation and a description of the simulations are given in \\S\\ref{sec:theory}. The resulting corrections to published measurements are discussed in \\S\\ref{sec:revise}. Our conclusions are summarised in \\S\\ref{sec:conc}. ", "conclusions": "" }, "1002/1002.4811_arXiv.txt": { "abstract": "We present results of new, deep \\emph{Suzaku} X-ray observations ($160$\\,ks) of the intracluster medium (ICM) in Abell~1689 out to its virial radius, combined with complementary data sets of the projected galaxy distribution obtained from the \\emph{SDSS} catalog and the projected mass distribution from our recent comprehensive weak and strong lensing analysis of \\emph{Subaru/Suprime-Cam} and \\emph{HST/ACS} observations. Faint X-ray emission from the ICM around the virial radius ($r_{\\rm vir}\\sim15\\farcm6$) is detected at $4.0 \\sigma$ significance, thanks to low and stable X-ray background of \\emph{Suzaku}. The \\emph{Suzaku} observations reveal anisotropic gas temperature and entropy distributions in cluster outskirts of $r_{500} \\simlt r \\simlt r_{\\rm vir}$ correlated with large-scale structure of galaxies in a photometric redshift slice around the cluster. The high temperature ($\\sim5.4~$keV) and entropy region in the northeastern (NE) outskirts is apparently connected to an overdense filamentary structure of galaxies outside the cluster. The gas temperature and entropy profiles in the NE direction are in good agreement, out to the virial radius, with that expected from a recent \\emph{XMM-Newton} statistical study and with an accretion shock heating model of the ICM, respectively. To the contrary, the other outskirt regions in contact with low density void environments have low gas temperatures ($\\sim 1.7$\\,keV) and entropies, deviating from hydrostatic equilibrium. These anisotropic ICM features associated with large-scale structure environments suggest that the thermalization of the ICM occurs faster along overdense filamentary structures than along low-density void regions. We find that the ICM density distribution is fairly isotropic, with a three-dimentional density slope of $-2.29\\pm0.18$ in the radial range of $r_{2500} \\simlt r \\simlt r_{500}$, and with $-1.24_{-0.56}^{+0.23}$ in $r_{500} \\simlt r \\simlt r_{\\rm vir}$, which however is significantly shallower than the Navarro, Frenk, \\& White universal matter density profile in the outskirts, $\\rho\\propto r^{-3}$. A joint X-ray and lensing analysis shows that the hydrostatic mass is lower than spherical lensing one ($\\sim60-90\\%$) but comparable to a triaxial halo mass within errors, at intermediate radii of $0.6r_{2500} \\simlt r \\simlt 0.8r_{500}$. On the other hand, the hydrostatic mass within $0.4r_{2500}$ is significantly biased as low as $\\simlt60\\%$, irrespective of mass models. The thermal gas pressure within $r_{500}$ is, at most, $\\sim50$--$60\\%$ of the total pressure to balance fully the gravity of the spherical lensing mass, and $\\sim30$--$40\\%$ around the virial radius. Although these constitute lower limits when one considers the possible halo triaxiality, these small relative contributions of thermal pressure would require additional sources of pressure, such as bulk and/or turbulent motions. ", "introduction": "Clusters of galaxies are the largest self-gravitating systems in the universe, and X-ray observations have revealed characteristics of intracluster medium (ICM) which carries $\\sim 70$\\% mass of known baryons, such as temperature, ICM mass, and metals therein. However, the ICM emission has been detectable only to $\\sim 0.6$ times the virial radius $r_{\\rm vir}$ \\citep[e.g.][]{pratt-2007} due to limited sensitivities of detectors. This means $\\sim 80$\\% of an entire cluster volume has been unexplored in X-rays. According to the hierarchical structure formation scenario based on cold dark matter (CDM) paradigm, mass accretion flows onto clusters, in particular, along filamentary structures, are still on-going. The ICM in the cluster outskirts, therefore, is expected to be sensitive to the structure formation and cluster evolution. However, since the X-ray data around the virial radius of most clusters were unavailable, we have not yet known the ICM characteristics in detail, such as density, temperature, pressure and entropy. The ICM around the virial radius is a fascinating frontier for X-ray observations. In the era of {\\it Suzaku} \\citep{mitsuda-2007}, detection of the ICM beyond 0.6 $r_{\\rm vir}$ has become possible thanks to low and stable particle background of the XIS detectors \\citep{koyama-2007}. In fact, \\citet{reiprich-2009} measured temperature profile of Abell~2044 out to $r_{\\rm 200}$, a radius within which the mean cluster mass density is 200 times the cosmic critical density, and found that the profile is in good agreement with hydrodynamic simulations. \\citet{george-2009} determined radial profiles of density, temperature, entropy, gas fraction, and mass up to $\\sim 2.5\\; h_{70}^{-1}$ Mpc of PKS~0745-191, and found that the temperature profile deviates from that expected from the ICM in hydrostatic equilibrium with a universal mass density profile as found by Navarro, Frenk \\& White (1996, 1997, hereafter NFW profile). \\citet{bautz-2009} also detected X-ray emission to $r_{\\rm 200}$ and found evidence for departure from hydrostatic equilibrium at radii as small as $r \\sim 1.3$ Mpc ($\\sim r_{\\rm 500}$). \\citet{fujita-2008} obtained metallicity of the ICM up to the cluster virial radii for the first time from a link region between the galaxy clusters Abell~399 and Abell~401. A gravitational lensing study is complementary to X-ray measurements, because lensing observables do not require any assumptions on the cluster dynamical states. The huge cluster mass distorts shapes of background galaxy images due to differential deflection of light paths, in a coherent pattern. As a result, multiple images, arcs and Einstein rings caused by the strong gravitational field of a cluster, appear around the cluster center, which is so-called strong gravitational lensing effect. It is capable of measuring mass distributions of a cluster central region very well. Outside the core, weak gravitational lensing analysis, which utilizes the coherent distortions in a statistical way, is a powerful tool to constrain the mass distribution out to the virial radius. In particular, the {\\it Subaru/Suprime-Cam} is the best instrument to conduct the weak lensing analysis \\citep[e.g.][]{miyazaki-2002,hamana-2003,broadhurst-2005b,umetsu-2008,umetsu-2009a,umetsu-2009b, hamana-2009,okabe-2008,okabe-2009,oguri-2009}, thanks to its high photon-collecting power and high imaging quality, combined with a wide field-of-view. Since the strong lensing analysis is complementary to weak lensing one, a joint analysis enables us to determine accurately the cluster mass profile from the center to the virial radius. Pioneer joint strong and weak lensing studies on Abell~1689 \\citep[][]{broadhurst-2005b,umetsu-2008} have shown that observed lensing profile is well fitted by an NFW mass profile with a high concentration parameter which is defined as a ratio of the virial radius to the scaled radius. Therefore, one of the best targets for the study of the ICM out to cluster outskirts is Abell~1689. Comparisons of X-ray observables with the well-determined lensing mass would allow us to conduct a powerful diagnostic of the ICM states, including a stringent test for hydrostatic equilibrium, for the first time. In addition, two mass measurements using X-ray and lensing analysis are of vital importance to understand the systematic measurement bias and construct well-calibrated cluster mass, which will improve the accuracy of determining cosmological parameters using a mass function of galaxy clusters \\cite[e.g. ][]{vikhlinin-2009,okabe-2010,zhang-2010}. In this paper, we describe the \\emph{Suzaku }observations and data reduction in \\S \\ref{section:obs}, and X-ray background estimation in \\S \\ref{section:xbkg}. In \\S \\ref{section:result}, we show results of spatial and spectral analyses and obtain radial profiles of temperature, electron density, and entropy. In \\S \\ref{section:dis}, we discuss the hydrostatic equilibrium around the virial radius, and compare X-ray observables, such as hydrostatic mass, pressure and entropy, with lensing mass derived from a joint strong and weak lensing analysis. We also discuss the ICM characteristics in the outskirts in light of its cosmological environment, such as the large-scale structure and low-density voids. We summarize the conclusions in \\S \\ref{section:summary}. In the present work, unless otherwise stated, we use $H_0 = 71\\; \\rm{km\\; s^{-1}\\; Mpc^{-1}}$ ($h = H_0 / 100\\; \\rm{km\\; s^{-1}\\; Mpc^{-1}} = 0.71$), assuming a flat universe with $\\Omega_{\\rm M} = 0.27$. This gives physical scale $1'' = 3.047$ kpc at the cluster redshift $z = 0.1832$ \\citep{struble-1999}. The definition of solar abundance is taken from Anders and Grevesse~(1988). Errors are given at the 90~\\% confidence level. ", "conclusions": "\\label{section:dis} \\subsection{Thermodynamics in the Outskirts} \\label{subsec:ICM_outskirts} In this section we discuss the distributions of gas temperature and entropy shown in figure \\ref{fig:prof}, particularly focusing on their anisotropic distributions in the last bin ($r_{500} \\simlt r \\simlt r_{\\rm vir}$). In the framework of the hierarchical clustering scenario based on CDM paradigm, mass aggregation processes onto clusters, in particular along the large-scale filamentary structures, are still on-going. Continuous mass accretion flows along filamentary structures, within which clusters are embedded, sometimes trigger {\\it internal shocks} in gas within clusters. It converts most of the tremendous amounts of kinetic energy into the thermal energy to heat the ICM. Therefore, the accretion flows are considered to play an important role in cluster evolution as well as in the ICM thermodynamics. If the shock heating dominates in the cluster outskirts, the entropy generated in the shock process becomes higher than those in the pre-shock regions. On the other hand, if the gas is adiabatically compressed during the infall with a sub-sonic velocity, the gas temperature increases but the entropy stays constant. According to analytical models and simulations considering both the continuous accretion shock and adiabatic compression processes \\citep[e.g.][]{tozzi-2001,voit-2002,borgani-2005}, the entropy profile is predicted to increase with $K\\propto r^{1.1}$ by the accretion shock, as long as the initial entropy is low, assuming that the ICM is in hydrostatic equilibrium and the kinetic energy of the infalling gas is instantly thermalized via shocks. This radial dependence has been generally confirmed in observed profiles \\cite[e.g. ][]{ponman-2003, pratt-2006} within $r_{500}$. Interestingly, our entropy profile in the Offset1 direction coincides very well with the model prediction up to the virial radius. It indicates a possibility that the ICM in the Offset1 $r_{500} \\simlt r \\simlt r_{{\\rm vir}}$ region is thermalized and close to the hydrostatic equilibrium. Indeed, the outskirts temperature in the Offset1 direction agrees with the predictions of the analytical models \\cite[e.g. ][]{komatsu-2001,tozzi-2001,ostriker-2005} in which, under the hydrostatic equilibrium assumption, the temperature around the virial radius declines at most $50$ percent of the inner values. Contrary to the Offset1, the sharp decline in both temperature and entropy in the Offset2-Offset4 regions indicates that the thermal pressure is not sufficient to balance the total gravity of the cluster. If the gas is not completely thermalized, the gas temperature should be lower than the value needed to maintain hydrostatic equilibrium. In this case, the accreted gas retains some fraction of the total energy in bulk motion, and the bulk pressure partly supports the gravity. In fact, recent numerical simulations \\cite[e.g.][]{vazza-2009} have shown that the kinetic energy of bulk motion carries $\\sim 30\\%$ of the total energy around the virial radius. In addition to the bulk pressure, pressure of turbulent motions in the ICM may also contribute to balance the total gravity \\cite[e.g. ][]{nagai-2007, piffaretti-2008, jeltema-2008, lau-2009,fang-2009,vazza-2009}. In order to study the validity for the assumption of hydrostatic equilibrium, we measured the hydrostatic equilibrium masses (H.E. mass) for the all, Offset1 and Offset234 regions. We here assume that the mass distribution is spherically symmetric and the ICM is in hydrostatic equilibrium at any radii. The hydrostatic mass, $M_{\\rm H.E.}$, is calculated from the parametric density and temperature profiles obtained by fitting data points in each azimuthal region (all, Offset1, \\& Offset234) with models (Appendix \\ref{appen:n+Tpara} and table \\ref{table:nfit}), as follows \\begin{eqnarray} \\frac{1}{\\rho_g} \\frac{d P_g}{d r} &=& - \\frac{GM}{r^2} \\label{eq:HE} \\\\ P_{g} &=& \\frac{\\mu_e}{\\mu} n_e k_B T, \\nonumber \\end{eqnarray} where $P_g$ is the gas pressure and $\\rho_g=\\mu_e m_p n_e$ is the gas mass density. The errors (68\\% CL uncertainty) are calculated with Monte Carlo simulations because model parameters are correlated with each other. The resulting mass is shown in figure~\\ref{fig:Mhe}. The cumulative mass in the Offset1 direction increases with radius, while the ones in all and Offset234 regions {\\it unphysically} decrease outside $\\sim7\\farcm0$--$8\\farcm0$. This {\\it unphysical} decrease in the cumulative mass is due to the sharp temperature drops in the Offset234 region. Here, in the parametric density model (equation \\ref{eq:ne_pr}), we adopted a modified $\\beta$ model to reproduce the density flatness in outskirts ($r>10\\farcm0$). The ratio of core radii $r_{c,2}/r_{c,1}$ in equation \\ref{eq:ne_pr} was fixed to be 3, since the flattening factor in which $r_{c,2}$ is included was not constraind well owing to few data points. However, even by choosing the ratio $\\Delta(r_{c,2}/r_{c,1})=\\pm1$, the radii, at which the masses have maximum values, are changed only by $+7\\%, -3\\%$ for Offset234 and $+12\\%,-6\\%$ for all regions, respectivly. This is because the curvature of the parametric density profile is insensitive to a choice of the ratio of core radii. Thus, the hydrostatic equilibrium we assumed is inadequate to describe the ICM in the outskirts except the Offset1 direction, that is, most of the ICM in the cluster outskirts is far from hydrostatic equilibrium. One alternative proposal to the bulk motion or turbulent motion as a cause of the non-hydrostatic equilibrium would be convective instability in the cluster outskirts. We investigated the possibility of convective instability in the radial direction. The convective motion in Offset234 is unstable outside $\\sim9\\farcm0$ but the time scale of the growing mode is comparable to the age of the universe at cluster redshift $z=0.1832$. Therefore, the convective instability is negligible for the {\\it unphysical} decrease in the cumulative mass. Another possible proposal is that ions have a higher temperature than that of electrons \\citep[]{fox-1997, ettori-1998, takizawa-1998}, and that the thermal pressure of ions supports the cluster mass. We computed the thermal equilibration time between electrons and ions through the Coulomb interaction. In the cluster outskirts of Abell~1689, the time scale is given by \\begin{eqnarray} t_{\\rm ei}\\sim 0.4 \\left(\\frac{n_e}{5\\times10^{-5}~{\\rm cm}^{-3}}\\right)^{-1} \\left(\\frac{k_B T}{2~{\\rm keV}}\\right)^{3/2}~{\\rm Gyr} \\end{eqnarray} \\cite[e.g. ][]{spitzer-1962,takizawa-1999,akahori-2009}. The Coulomb interaction would provide us with the maximum time scale of thermal equilibrium, because there might be other processes, such as plasma instability, to facilitate the interaction between electrons and ions. Since the resulting time is much shorter than the dynamical time scale, it is difficult to explain the {\\it unphysical} decrease. Therefore, the bulk and/or turbulent motions would be the main source(s) of additional pressure needed in the Offset234 directions. In \\S \\ref{subsection:Pwl}, we shall discuss on this point again based on a joint X-ray and lensing analysis. What makes the anisotropic distributions of gas temperature and entropy in the outskirts? The internal shocks associated with mass accretion flows would heat the ICM. Numerical simulations \\citep[e.g.][]{ryu-2003, molnar-2009} have shown that, the low-density gas in low-density regions, so-called voids, accretes onto the vicinity of clusters by the gravitational pull, with an order of thousands of kilometers per second. As a result, the internal shocks form outside virial radius at which gas density sharply changes. It is also referred to as {\\it virial shocks} \\citep{ryu-2003}. Meanwhile, the accretion flow along filamentary structures forms internal shocks inside the virial radius. These shocks migrate relatively deep into a cluster potential well, because a large amount of matter, such as gas, galaxies and groups, accretes from filaments \\citep{ryu-2003}. Therefore, the thermalization in the outskirts along the direction of filamentary structure takes place faster. Based on this accretion shock scenario, the observed anisotropic distributions of gas temperature and entropy in the outskirts inside virial radius would be associated with large-scale structures. Another possibility for the anisotropy is that infalling galaxies, especially massive galaxies, might be reservoirs of cold gas. The cold and dense gas, which is bound by the dark matter potential of galaxies, would decrease the emission-weighted temperature in the outskirts. Although we excluded point sources in the spectral study, we cannot rule out this possibility only from the X-ray results. If this is the case, member galaxies around the virial radius ($10\\farcm0$--$18\\farcm0$) in the Offset234 directions would be more numerous than those in the Offset1 direction. In order to search the large-scale structures or the anisotropic galaxy distribution around the virial radius, we made maps of galaxy number density using the {\\it SDSS} DR7 catalogue (Abazajian et al. 2009) from {\\it SDSS} CasJobs site\\footnote{http://casjobs.sdss.org/}. We first looked into bright red-sequence galaxies which are expected to host cold gas. We selected red-sequence galaxies by criteria of $|(g'-i') - (-0.048i'+2.556)|<0.152$ and $i'<21~{\\rm ABmag}$, where we used {\\it psfMag} for magnitude and {\\it modelcolor} for color. We made projected galaxy distribution, smoothed with a Gaussian kernel, $w=\\exp(-r^2/r_g^2)$, of FWHM$=2(\\ln2)^{1/2}r_g=1\\farcm67$. The resulting map is shown in the left panel of figure \\ref{fig:Mapsdss}. We also looked into the same map derived from {\\it Subaru} data. It is consistent with the {\\it SDSS} map, although a part of the outer region is not available in the {\\it Subaru} data due to the limitation of its field-of-view. In figure \\ref{fig:Mapsdss} (left), the central part of density map is clearly elongated in the north-south direction. \\citet{umetsu-2008} have found the same elongation in a two-dimensional mass distribution. However, we could find neither an apparent feature nor a significant difference between the projected distributions in the Offset1 and the other directions in the cluster outskirts of $10\\farcm0 < r < 18\\farcm0$. Therefore, the temperature and entropy anisotropies in the outskirts are not due to the cold gas in galaxies. Next, we made a larger map to investigate the presence of large-scale structures associated with Abell~1689. \\cite{lemze-2009} has shown that the maximum line-of-sight velocity of infall galaxies is $\\sim4000~{\\rm km/s}$. Therefore, taking into account the uncertainty of redshift due to line-of-sight velocities, we selected galaxies in the range of $|z-z_{\\rm c}| < \\delta z =\\sigma_{v,{\\rm max}}(1+z_{\\rm c})/c\\simeq0.0158$, where $z_{\\rm c}$ is the cluster redshift, $z$ is a photometric redshift, $\\sigma_{v,{\\rm max}}=4000~{\\rm km/s}$, and $c$ is the light velocity. The mean photometric redshift for our sample is $\\bar{z}=0.18320\\pm0.00004\\pm0.00035$, where the first error is the statistical error and the second one is the systematic error due to photometric uncertainty of each galaxy. Our sample of galaxies thus statistically represents a slice around Abell~1689 in redshift space. The resulting map of galaxies smoothed with a Gaussian kernel of FWHM=20\\farcm0 is shown in the right panel of figure \\ref{fig:Mapsdss}. In the Offset1 direction, a filamentary overdensity region outside the virial radius is found, where the galaxy number density is more than 1.5 times higher than those in the other directions. In the Offset4 direction outside the FOV of XIS, a narrow sheet of galaxies is also found, which is connected to a northwest overdensity region. We note that the projected elongated direction of cluster mass and galaxies (left panel of figure \\ref{fig:Mapsdss}) does not coincide with the filamentary direction (right panel), suggesting that a mass structure seen in the central region of a cluster is not necessarily connected directly to a filament. Our result does not change even when we choose twice or half $\\delta z$ in the galaxy selection. We also tried to investigate the line-of-sight filamentary structure by a Monte -Carlo simulation computing the redshift distribution of galaxies with photometric errors. However, we could not obtain reliable results within a redshift resolution smaller than $\\sim10~{\\rm Mpc}$. We could not see the line-of-sight structure from the spectral data of {\\it SDSS} either, because available number of galaxies is too small. The high temperature and entropy region in the outskirts is clearly correlated with the galaxy overdensity region associated with the large-scale structure outside Abell~1689, as shown in figure~\\ref{fig:Map_kt_sdss}. This indicates that the ICM in the outskirts is significantly affected by surrounding environments of galaxy clusters, such as the filamentary structures and the low density void regions. The large-scale structure would play an important role in the thermalization process of the ICM in the outskirts. In particular, our result suggests that the thermalization in the outskirts along the filamentary structure takes place faster than that in the void region. \\subsection{A Joint X-ray and Lensing Analysis} \\label{sec:joint} In this subsection, we carry out a joint X-ray and lensing analysis, incorporating {\\it Subaru/Suprime-Cam} and {\\it HST/ACS} data. Abell~1689 has been the focus of intensive lensing studies in recent years \\citep[e.g, ][]{broadhurst-2005a,broadhurst-2005b,limousin-2007,umetsu-2008,corless-2009}. It has been shown by \\cite{broadhurst-2005b} and \\cite{umetsu-2008} that joint lensing profiles of Abell~1689 obtained from their ACS and Subaru data with sufficient quality are consistent with a continuously steepening density profile over a wide range of radii, $r=10-2000$\\,kpc$\\,h^{-1}$, well described by the general NFW profile \\citep{navrro-1996,navrro-1997}, whereas the singular isothermal sphere model is strongly disfavored. They also revealed that the concentration parameter of the NFW profile, namely the ratio of the virial radius to the scale radius, $c_{\\rm vir}=r_{\\rm vir}/r_s$, is much higher than predicted by cosmological $N$-body simulations \\cite[e.g., ][]{bullock-2001,neto-2007}. We note that the virial radius, $r_{\\rm vir}$, within which the mean interior density is $\\Delta_{\\rm vir}\\simeq110$ \\citep{nakamura-1997} times the critical mass density, $\\rho_{\\rm cr} (z)$, at a cluster redshift is given by \\begin{eqnarray} M_{\\rm vir}= \\frac{4}{3} \\pi \\Delta_{\\rm vir} \\rho_{\\rm cr}(z) r_{\\rm vir}^3. \\label{eq:viral} \\end{eqnarray} In the present paper, the full lensing constraints derived from the joint ACS and Subaru data allow us to compare, for the first time, X-ray observations with the total mass profile for the entire cluster, from the cluster center to the virial radius. The lensing analysis is free from any assumptions about the dynamical state of the cluster, so that a joint X-ray and lensing analysis provides a powerful diagnostic of the ICM state and any systematic offsets between the two mass determination methods. Here we first summarize our lensing work on Abell~1689 before presenting the results from our joint analysis. \\cite{umetsu-2008} combined HST/ACS strong lensing data with Subaru weak lensing distortion and magnification data in a two-dimensional analysis to reconstruct the projected mass profile. Their full lensing method, assuming the spherical symmetry, yields best-fit NFW model parameters of $M_{\\rm vir}=1.47^{+0.59}_{-0.33} \\times 10^{15} M_\\odot\\,h^{-1}$ and $c_{\\rm vir}=12.7\\pm 2.9$ (including both statistical and systematic uncertainties; see also \\cite{lemze-2009}), which properly reproduce the observed Einstein radius of $\\theta_{\\rm E}=45\\arcsec$ for $z_s=1$ \\citep{broadhurst-2005a}. Here we deproject the two-dimensional mass profile of \\cite{umetsu-2008} and obtain a non-parametric $M_{\\rm 3D}$ profile simply assuming spherical symmetry \\citep{broadhurst-2008,umetsu-2009b}. This method is based on the fact that the surface-mass density $\\Sigma_m(R)$ is related to the three-dimensional mass density $\\rho(r)$ by an Abel integral transform; or equivalently, one finds that the three-dimensional mass $M_{\\rm 3D}($28 mag/arcsec, they are thus named ``ultra-faint\" dSphs. They have properties intermediate between globular clusters (GCs) and dSphs, and contain very metal poor stars, as metal poor as [Fe/H]$\\sim -3.0, -4.0$ dex. Often the ``ultra-faint\" dSphs have irregular shape and appear to be distorted by the tidal interaction with the MW. Like their ``bright\" counterparts, they have high mass to light ratios. All ``ultra-faint\" dSphs host an ancient population as old as $\\sim$ 10 Gyr, and have GC-like CMDs, resembling the CMDs of metal poor Galactic clusters like M92, M15 and M68. The CMD of an ``ultra-faint\" dSph galaxy, Ursa Major~I (UMa~I), is shown in Fig. 2. \\begin{figure}[!h] \\centerline{\\hbox{\\psfig{figure=clementini-fig2.eps,angle=0,clip=,width=9cm}}} \\caption[]{The $V, B-V$ color magnitude diagram of the UMa~I ``ultra-faint\" dSph galaxy, from Garofalo (2009). The solid line is the ridge line of the Galactic globular cluster M68 which has been shifted in magnitude and color to match the horizontal branch and the red giant branch of UMa~I. Solid circles mark the RR Lyrae stars detected in the galaxy.} \\label{clementini-fig2} \\end{figure} The number of the M31 satellites is poorly constrained. Twelve dwarf galaxies were known to be M31 companions until 2004, among which only 6 dSphs, but several new M31 satellites were discovered afterwards. The most recent census is reported by Martin et al. (2009) along with the discovery of two new M31 dSph satellites: And XXI and And XXII. As for the MW, the number of M31 satellites is likely to increase significantly in the near future. Fig. 3 shows the location of ``bright\" and ``ultra-faint\" dSphs in the absolute magnitude versus half-light radius ($r_h$) diagram. The plot is an adapted and updated version of Belokurov et al. (2007) Fig. 8. The MW globulars and some of the M31 GCs are also shown in the figure, for a comparison. With their faint luminosities and large dimensions, the ``ultra-faint\" dwarfs sample a totally unexplored region of $M_V - \\log (r_h)$ plane. Among the ``ultra-faint\" dSphs only Canes Venatici~I (CVn~I), the brightest of these systems, lies close to the ``bright\" dSphs. \\begin{figure}[!h] \\centerline{\\hbox{\\psfig{figure=clementini-fig3_new2.eps,angle=0,clip=,width=11.5cm}}} \\caption[]{Location of ``bright\" (``old dSphs\" in the figure) and ``ultra-faint\" dSphs around the MW and M31 in the absolute magnitude versus half-light radius ($r_h$) diagram. Adapted and updated from Belokurov et al. (2007).} \\label{clementini-fig3} \\end{figure} ", "conclusions": "Most of the RRab stars in the MW and M31 halos have Oosterhoff type I properties, but the inner MW halo also contains a significant Oo~II component. By contrast the Oo~I type is rare among the dSph galaxies, where only Sagittarius is Oo~I. The ``bright\" MW dSphs tend to be Oo-Intermediate, while the ``ultra-faint\" dSphs tend to have Oo~II type, in so far they can be classified. The Galaxy is unlike to have formed by accretion of protogalactic fragments resembling its present-day ``bright\" dwarf satellites. On the other hand, systems resembling the ``ultra-faint\" dSphs may have contributed in the past to the formation of the MW halo. The study of the M31 RR Lyrae population is still in pioneering stage. Further data for both field and cluster variables are needed to reach firm conclusions on the Oosterhoff classification of the Andromeda galaxy." }, "1002/1002.4610.txt": { "abstract": "A version of the Swiss-cheese model is investigated. The flat Friedmann-Robertson-Walker (FRW) universe is modified by the addition of several spherical regions with Lema\\^itre-Tolman-Bondi metric. We discuss light propagation in this model in detail to pave the way for a detailed numerical study of the Hubble diagram. %%We discuss this model in detail, and show that such a model universe can have %%a realistic density distribution on large scales, and exhibits accelerating %%expansion for a limited period of time. Especially, we determine the Hubble %%diagram and discuss its properties. ", "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Type Ia supernova data interpreted in terms of homogeneous, isotropic cosmological models necessitate the introduction of a huge amount of dark energy, i.e., a substance of negative pressure (like the cosmological constant). A recent alternative suggestion is the consideration of inhomogenity in the density distribution. Indeed, lensing effects of local matter abundances modify the positions of distant objects on the Hubble diagram. There are several recent papers about this question without a concordance whether this effect is strong enough to give account of supernova data,[] albeit the majority of the authors seems to deny this possibility. The question is rather important since the existence of dark energy (or cosmological constant) profoundly modifies our knowledge about matter, already challenged by the existence of dark matter. In the present paper we consider the problem within the framework of the exactly solvable Swiss-cheese model. The version of the model we use is constructed the following way. Nonoverlapping spheres are cut from a flat Friedmann-Robertson-Walker (FRW) universe. The mass they contained before is compressed within each sphere to a smaller sphere with homogeneous density distribution. Hence the inner spheres form sections of some closed FRW model. Between the outer and inner spheres there is a vacuum, where, due to spherical symmetry, the Schwarzschield metric describes the gravitational field. Within the inner spheres the closed FRW metric is valid, while outside the cut spheres the flat FRW metric is relevant. The metric and its first derivatives are continuous across the bordering surfaces of the different regions. We use the Landau conventions, i.e., we assume $+---$ signature for the metric. The zeroth component of a four-vector is timelike, the first, second and third components are spacelike. Four-vectors are indexed by Latin letters, three-vectors by Greek letters. Throughout we use $c=1$ units. At light propagation we use the index $1$ for labeling initial quantities and the index $0$ for labeling final (i.e., present) quantities. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ", "conclusions": "" }, "1002/1002.2349_arXiv.txt": { "abstract": "We perform fully-kinetic particle-in-cell simulations of an hot plasma that expands radially in a cylindrical geometry. The aim of the paper is to study the consequent development of the electron temperature anisotropy in an expanding plasma flow as found in a collisionless stellar wind. Kinetic plasma theory and simulations have shown that the electron temperature anisotropy is controlled by fluctuations driven by electromagnetic kinetic instabilities. In this study the temperature anisotropy is driven self-consistently by the expansion. While the expansion favors an increase of parallel anisotropy ($T_\\parallel>T_\\perp$), the onset of the firehose instability will tend to decrease it. We show the results for a supersonic, subsonic, and static expansion flows, and suggest possible applications of the results for the solar wind and other stellar winds. ", "introduction": "The steady-state solution of the most simple dynamical equations that describe the evolution of a stellar wind dates back to the hydrodynamic model for the solar wind proposed by Parker in 1958 \\citep{parker58}. Increasingly sophisticated models have been developed to take into account the different observed characteristics of the solar wind, and modeling has been based on analytic, semi-analytic and simulational methods \\citep{hollweg08}. The magnetohydrodynamics (MHD) approach is widely used for global modeling, but since it treats the plasma as a fluid it does not include any effects due to non-Maxwellian particle velocity distributions. More recently, kinetic simulations have tried to take in account and explain features related to the observed non-thermal distribution of particles in velocity space (see e.g. \\citet{landi2003, zouganelis05} ). The aim of such exospheric models is to provide a realistic description of the wind dynamics that includes the transition from a collision-dominated to a collisionless region. In doing so, however these models do not include the effect of plasma instabilities and therefore cannot be regarded as completely self-consistent. For the solar wind, the importance of small scale fluctuations, associated with kinetic plasma instabilities generated by non-Maxwellian particle distributions, is now widely recognized. This has come about through a convergence of observations, theory and simulations. It is argued that many macroscopic quantities that characterize the solar wind, such as the particle temperature anisotropy or the electron heat flux, are always observed with values that are bounded by the possible onset of kinetic instabilities. The general argument is the following: the expansion of the flow leads to a distortion of the distribution function, which represents, for example, an increase in the temperature anisotropy. The distortion of the distribution function can be enough to trigger linear instability, i.e., there is some free-energy available which can create plasma fluctuations. Based on simulations of the initial value problem for such instabilities, it is evident that the fluctuations act to reduce the distortion of the distribution function. In other words, the primary effect of the fluctuations produced in the instability is to restore a stable (or marginally stable) distribution function. In the case of a temperature anisotropy driven by expansion, we might expect the onset of instabilities associated with temperature anisotropy to act as an upper limit for the anisotropy. This argument can be broadly applied to many situations, such as formation of heat flux, or indeed compression flows. In the solar wind context it is expected that the expansion of the wind away from the Sun produces a parallel temperature anisotropy $T_\\parallel>T_\\perp$ (with reference to the magnetic field direction) due to the conservation of adiabatic invariants. If one adopts a fluid viewpoint (for instance the `double-adiabatic' theory \\citep{CGL} that is often used in this context), the anisotropy is not bounded and it should increase indefinitely for increasing distance from the Sun. However it has been observed that the temperature anisotropy is limited in value both for ions \\citep{gary01,kasper02,hellinger06,matteini07} and electrons \\citep{gary05,stverak08}, and it has been suggested that this is due to the onset of kinetic instabilities, such as the firehose for $T_\\parallel>T_\\perp$ and the mirror or ion-cyclotron instabilities for $T_\\parallelN^2$. We find this is no longer true when the medium is a magnetized plasma. We investigate the effect of stable stratification on magnetic field and velocity profiles unstable to magneto-shear instabilities, i.e., instabilities which require the presence of both magnetic field and shear flow. We show that a family of profiles originally studied by \\citet{td06} remain unstable even when $V'^2/4V'^2/4 $, we always found stability. We provided two heuristics for understanding the destabilization due to magnetic fields. An inhomogeneous magnetic field provides a free energy source which can be tapped by an instability. Thus, while a homogeneous magnetic field can be destabilizing because vorticity is no longer frozen into the flow, allowing new unstable plasma motions, only an inhomogeneous field can provide the source of energy needed to violate Richardson's criterion. We briefly applied our results to the solar tachocline and to high mass, rapidly rotating stars. In the bulk of the tachocline, $\\mbox{Ri}$ is very large because the Sun rotates slowly. Very near the boundary of the convection zone, $\\mbox{Ri}$ drops because $N^2$ is passing through zero. A similar situation holds, for different reason, in high mass stars. Although these stars rotate rapidly, the regions of strong shear coincide with regions of strong, stabilizing, molecular weight gradient. This keeps $\\mbox{Ri}$ large, except near convection zone boundaries. Thus, in stars, the destabilization of stratified shear flow by magnetic fields is most likely to occur in thin regions on the stable side of convection zone boundaries. If the weakening of buoyancy by thermal diffusion destabilizes magnetized flow in the same way as unmagnetized flow, the unstable region could be much larger, however. Our 2D slab model is not a realistic geometry for many applications. The introduction of additional terms, such as curvature terms from toroidal geometry or the centrifugal force for rotation, probably changes our results quantitatively, but not qualitatively. The Boussinesq approximation could also be relaxed to allow more realistic density profiles and other physics. Inclusion of diffusive effects would allow us to consider non-ideal instabilities, including the secular shear instability. For many applications, the non-linear phase and saturation of these instabilities is also important for determining effects such as angular momentum transport. These considerations should be investigated further to better understand the nature of magneto-shear instabilities in a stratified medium." }, "1002/1002.1112_arXiv.txt": { "abstract": "To exploit synergies between the {\\em Herschel} Space Observatory and next generation radio facilities, we have extended the semi-empirical extragalactic radio continuum simulation of Wilman et al. to the mid- and far-infrared. Here we describe the assignment of infrared spectral energy distributions (SEDs) to the star-forming galaxies and active galactic nuclei, using {\\em Spitzer} 24, 70 and 160\\micron~and {\\em SCUBA} 850\\micron~survey results as the main constraints. Star-forming galaxies dominate the source counts, and a model in which their far-infrared--radio correlation and infrared SED assignment procedure are invariant with redshift underpredicts the observed 24 and 70\\micron~source counts. The 70\\micron~deficit can be eliminated if the star-forming galaxies undergo stronger luminosity evolution than originally assumed for the radio simulation, a requirement which may be partially ascribed to known non-linearity in the far-infrared--radio correlation at low luminosity if it evolves with redshift. At 24\\micron, the shortfall is reduced if the star-forming galaxies develop SEDs with cooler dust and correspondingly stronger Polycyclic Aromatic Hydrocarbon (PAH) emission features with increasing redshift at a given far-infrared luminosity, but this trend may reverse at $z>1$ in order not to overproduce the sub-mm source counts. The resulting model compares favourably with recent {\\em BLAST} results and we have extended the simulation database to aid the interpretation of {\\em Herschel} surveys. Such comparisons may also facilitate further model refinement and revised predictions for the {\\em Square Kilometer Array} and its precursors. ", "introduction": "In Wilman et al.~(2008) (hereafter W08) we presented a semi-empirical simulation of the extragalactic radio continuum sky primarily intended to aid the design of scientific programmes for the next generation of radio telescope facilities, culminating in the {\\em Square Kilometer Array} (SKA). The simulation covers a sky area of $20 \\times 20$~deg$^{2}$ and contains radio-loud and radio-quiet active galactic nuclei (AGN) and star-forming galaxies out to redshift $z=20$ within a framework for their large-scale clustering. The full source catalogue -- containing 320 million sources above 10~nJy at five frequencies ranging from 151~MHz to 18~GHz -- can be accessed at the SKADS Simulated Skies ($S^{3}$) database\\footnote{http://s-cubed.physics.ox.ac.uk} under $S^{3}$--SEX (semi-empirical extragalactic simulation). There are by necessity numerous uncertainties and limitations in the $S^{3}$--SEX simulation, including but not limited to issues such as the form of the high-redshift evolution of the AGN and galaxies, the lack of star-forming/AGN hybrid galaxies, and the abundance of highly-obscured Compton-thick AGN. In so far as possible, flexibility was built into the simulations to allow post-processing to improve their accuracy as observations in the years ahead lead to improved constraints. Major advances in these areas are expected from the far-infrared surveys to be performed by the {\\em Herschel} Space Observatory (Pilbratt 2008). To facilitate such comparisons we have post-processed the $S^{3}$--SEX simulation to cover these wavelengths. In this paper, we describe the recipes for assigning infrared spectral energy distributions (SEDs) to the radio sources, using existing mid- and far-infrared results from {\\em Spitzer} and sub-mm survey data from {\\em SCUBA} as the primary constraints. We then present our predictions for {\\em Herschel} surveys, bolstered by a comparison with results from the Balloon-borne Large Aperture Submillimetre Telescope ({\\em BLAST}) which offer a foretaste of {\\em Herschel}'s capabilities. In keeping with the philosophy of the S$^{3}$ project to maximise the degree of interaction between the user and the database, the infrared fluxes are also provided on the $S^{3}$ webpage from which users can generate simulation products for comparison with observations. The radio--infrared connection is of immense empirical and theoretical interest for extragalactic surveys, for the identification and follow-up of {\\em Herschel} sources, and for probing the physics of the far-infrared--radio correlation and its possible evolution. The simulation can nevertheless also serve as a standalone {\\em Herschel} simulation, for comparison with others such as the phenomenologically-inspired backward evolution models of Valiante et al.~(2009) and Pearson \\& Khan~(2009), and the model of Lacey et al.~(2009) which is based on a semi-analytical galaxy formation model. ", "conclusions": "We have post-processed the S$^{3}$-SEX semi-empirical simulation of the extragalactic radio continuum sky (W08) to make predictions for {\\em Herschel} surveys at far-infrared wavelengths. Existing observations in the mid-infrared with {\\em Spitzer} at 24, 70 and 160\\micron~and at 850\\micron~with {\\em SCUBA}, together impose strong constraints on the assignment of infrared SEDs to the star-forming galaxies as a function of redshift. Our principal findings, incorporated into the final model, are as follows: (i) In order to match the 70\\micron~counts, the star-forming galaxies are required to undergo stronger luminosity evolution than assumed by W08 for the radio simulation. It may be that W08 simply used an inaccurate evolutionary prescription based on the information available at that time; alternatively, it could be that there is genuine differential evolution between the far-infrared and radio populations, as a result of an evolving non-linearity in the far-infrared:radio correlation which is already apparent locally at low luminosity . Results from {\\em Herschel} and {\\em SKA} precursors will resolve this issue. (ii) From the local Universe to redshift $z=1$, star-forming galaxies are required to develop progressively cooler SED templates at fixed L(FIR) once the latter has been set by the far-infrared--radio correlation; the rest-frame 60--100\\micron~colour is assumed to evolve as $\\rm{log} L_{\\rm{60}}/L_{\\rm{100}} \\sim (1+z)^{-2.5}$; beyond $z=1$, this evolution must go into reverse in order not to overproduce the 850\\micron~source counts. Our chosen model compares favourably with recent far-infrared survey results from {\\em BLAST}. This inspires confidence in using it to make predictions for {\\em Herschel} Key Programme surveys. Our predicted source counts and redshift distributions correspond closely with those of Valiante et al.~(2009) and Lacey et al.~(2009), respectively, despite the clear differences in the methodologies of our models. This suggests that the combination of {\\em Spitzer} mid-infrared and 850\\micron~data impose strong constraints on the allowable model parameter space. Data products from our simulation are available on the S$^{3}$ website\\footnote{http://s-cubed.physics.ox.ac.uk} and we expect them to serve as a valuable resource for the interpretation of {\\em Herschel} surveys." }, "1002/1002.1865_arXiv.txt": { "abstract": "{We report the results of our \\xmm\\ monitoring of \\sn. The ongoing propagation of the supernova blast wave through the inner circumstellar ring caused a drastic increase in X-ray luminosity during the last years, enabling detailed high resolution X-ray spectroscopy with the Reflection Grating Spectrometer.} {The observations can be used to follow the detailed evolution of the arising supernova remnant.} {The fluxes and broadening of the numerous emission lines seen in the dispersed spectra provide information on the evolution of the X-ray emitting plasma and its dynamics. These were analyzed in combination with the EPIC-pn spectra, which allow a precise determination of the higher temperature plasma. We modeled individual emission lines and fitted plasma emission models.} {Especially from the observations between 2003 and 2007 we can see a significant evolution of the plasma parameters and a deceleration of the radial velocity of the lower temperature plasma regions. We found an indication (3$\\sigma$-level) of an iron K feature in the co-added EPIC-pn spectra.} {The comparison with Chandra grating observations in 2004 yields a clear temporal coherence of the spectral evolution and the sudden deceleration of the expansion velocity seen in X-ray images $\\sim$6100 days after the explosion.} ", "introduction": "The circumstellar ring system around \\SN\\ was ejected by the progenitor star approximately 20\\,000 years before the supernova explosion. About 10 years after the explosion the blast wave started to propagate through the inner ring, causing compression, heating and ionization of its matter. At that time bright unresolved regions, so called hot spots, appeared all around the ring in HST images \\citep{2000ApJ...537L.123L} and the rather linear increase in the soft X-ray band, as observed with ROSAT since 1992 \\citep{1994A&A...281L..45B,1996A&A...312L...9H}, was followed by an exponential brightening as monitored with \\xmm, Chandra, Suzaku and Swift. Recently, some flattening of the flux increase is seen \\citep[][and references therein for the individual fluxes]{2009PASJ...61..895S}. The X-ray spectra are interpreted as thermal emission composed of a lower temperature component ($\\sim$0.5~keV) and a higher temperature component ($\\sim$2.5~keV). Deep Chandra grating observations allowed the measuring of the bulk gas velocity due to spectral line deformation. Surprisingly these values were lower than the velocities expected from the plasma temperatures. This suggests the contribution of reflected shocks, additional to the 'normal' forward shock \\citep{2005ApJ...628L.127Z,2006ApJ...645..293Z,2009ApJ...692.1190Z,2008ApJ...676L.131D}. The expansion velocity derived from Chandra X-ray images is even higher. About 6100 days after the explosion a sharp deceleration to 1600$-$2000~\\kms\\ was observed \\citep{2009ApJ...703.1752R}. Shocks transmitted into denser regions of the ring have a slower shock wave velocity and therefore can be responsible for the low temperature component. This interpretation is supported by the morphology seen in optical and X-ray images \\citep{2007AIPC..937....3M}, as well as the similar evolution of the soft X-ray flux (0.5-2.0~keV) and the evolution of highly ionized optical emission lines from the hot spots \\citep{2006A&A...456..581G}. This emission might be caused by even slower radiative shocks. Renewed radio emission of \\SN\\ was detected in July 1990 \\citep{1990IAUC.5086....2T}. A continuously rising flux and increasing source radius has been observed by the ATCA since then \\citep{2007AIPC..937...86G,2008ApJ...684..481N}. The increase of the radio light curve matches the evolution of the hard X-ray flux (2-10~keV) quite closely \\citep{2005ApJ...634L..73P,2007AIPC..937...33A}, thus the synchrotron radio emission may originate in the hot thermal plasma between the forward and reverse shock and a fraction of the hard X-ray flux may also have a non-thermal origin. Broad \\Halp\\ and Ly$\\alpha$ lines suggest the presence of the reverse shocks \\citep{2003ApJ...593..809M,2006ApJ...644..959H}. This study reports on the yearly \\xmm\\ monitoring observations of \\SN\\ between January 2007 and January 2009 together with one prior observation from May 2003, making it possible to reveal the detailed plasma evolution. We measured light curves and widths of individual emission lines and analyzed the spectral plasma evolution by fitting thermal plasma emission models. ", "conclusions": "\\begin{enumerate} \\item With our monitoring we can follow the detailed evolution of the supernova remnant of \\SN. E.g. the upturn in ionization age and emission measure ratio between 2003 and 2007 shows that the blast wave was propagating into the inner ring. \\item The decreasing line widths at lower energies in between the first two observations indicate a deceleration of the lower temperature plasma, that correlates well with the decelerating ring expansion as observed in the Chandra images. \\item The electron temperature derived for the soft temperature component with plasma models is consistent with the line widths of the emission lines in the corresponding energy range. This is expected, if the emission is primary caused by shocks transmitted into denser regions. \\item The lower statistics at higher energies do not justify such conclusions for the high temperature component, where the bulk velocity can be reduced by the contribution of reflected shocks which would also heat the plasma further. But the decreasing temperature with time and the higher line widths seen by \\xmm\\ and Chandra rather suggest forward shocks in less denser regions as the dominating process. \\item The iron K feature argues rather for a thermal high energy component and little if any non-thermal contribution. The line shape further suggests the contribution ($\\sim$50\\%) of a cold iron line (2.3$\\sigma$ level), possibly emitted from the supernova debris. \\end{enumerate} Further monitoring will yield information on the arising SNR of \\SN\\ and the increasing statistics of the brightening source will allow even more detailed analyses also beyond the assumption of plane-parallel geometry. This might put further constraints on the origin of the X-ray emitting regions and their dynamics. Similarly the iron K emission can yield information on the evolution of the debris." }, "1002/1002.4447_arXiv.txt": { "abstract": "{GJ$\\:$436b is the first extrasolar planet discovered that resembles Neptune in mass and radius. Two more are known (HAT-P-11b and Kepler-4b), and many more are expected to be found in the upcoming years. The particularly interesting property of Neptune-sized planets is that their mass $M_p$ and radius $R_p$ are close to theoretical $M$-$R$ relations of water planets. Given $M_p$, $R_p$, and equilibrium temperature, however, various internal compositions are possible.} {A broad set of interior structure models is presented here that illustrates the dependence of internal composition and possible phases of water occurring in presumably water-rich planets, such as GJ$\\:$436b on the uncertainty in atmospheric temperature profile and mean density. We show how the set of solutions can be narrowed down if theoretical constraints from formation and model atmospheres are applied or potentially observational constraints for the atmospheric metallicity $Z_1$ and the tidal Love number $k_2$.} {We model the interior by assuming either three layers (hydrogen-helium envelope, water layer, rock core) or two layers (H/He/H$_2$O envelope, rocky core). For water, we use the equation of state H$_2$O-REOS based on finite temperature - density functional theory - molecular dynamics (FT-DFT-MD) simulations.} {Some admixture of H/He appears mandatory for explaining the measured radius. For the warmest considered models, the H/He mass fraction can reduce to $10^{-3}$, still extending over $\\sim 0.7\\RE$. If water occurs, it will be essentially in the plasma phase or in the superionic phase, but not in an ice phase. Metal-free envelope models have $0.021.2$) and highly flattened halos ($q<0.6$) are strongly excluded if our approximation of a true isothermal halo is valid. Our mass modelling analysis assumed that the \\HI\\ gas velocity dispersion was vertically isothermal, as no measurements of the vertical variation of the \\HI\\ gas velocity dispersion are as yet available.\\footnote{Our measurements of the \\HI\\ gas velocity dispersion used the vertically averaged \\HI\\ distribution, ie. they are the luminosity-weighted mean velocity dispersion as a function of radius.} If the \\HI\\ velocity dispersion is in fact vertically declining, this would lead to a larger estimated vertical gradient of the total vertical force, which may allow a pseudo-isothermal model for the halo. UGC7321 is a gas-rich galaxy ($M_{\\HI}/L_R=2.2$), with a very low stellar mass galaxy ($M=3\\times10^8$ \\msun), four times less massive than the gas disk. The $R$-band stellar mass-to-light ratio of UGC7321 is very low at $M/L_R\\lesim0.2$. Mass modelling of the vertical force distribution showed that vertical force fitting provides a much stronger constraint on the stellar mass-to-light ratio than the standard method of radial force fitting via rotation curve decomposition. Two important assumptions in this work need to be tested further. The first is that the \\HI\\ velocity dispersion is isothermal in $z$. For a definitive estimate of the $R$ and $z$-components of the total potential gradients from \\HI\\ hydrostatics, it is essential to have reliable measurements of the \\HI\\ density, rotation and velocity dispersion as a function of both $R$ and $z$. It should be possible, with additional short spacing ATCA observations supplementing our data, to measure the \\HI\\ velocity dispersion as a function of $z$ in ESO274-G001 by modelling the \\HI\\ XV diagram at varying heights above the galactic plane. ESO274-G001 is the closest, isolated, southern edge-on galaxy at a distance of 3.4 Mpc. In the northern hemisphere, UGC7321 is a prime candidate due to its high \\HI\\ mass, despite its larger distance of 10 Mpc. The large \\HI\\ flaring means that the \\HI\\ could be measured at a height of 400 pc for radii from 5-11 kpc, and at 700 pc for radii from 9-11 kpc. The dwarf Scd galaxy NGC5023 is also an excellent candidate, given its distance of about 8 Mpc \\citep{vdks1982a}. For this galaxy, early flaring measurements by \\citet{bsvdk1986} found that the gas thickness was constant with radius. This is a surprising result, because the $v_{max}$ value for this galaxy is only about $80$ \\kms, and large flaring might be expected. It would be interesting to measure the radial and vertical variation of the gas velocity dispersion, gas flaring, and halo shape using beter data, as this galaxy has a similar size, \\HI\\ brightness and total mass as UGC7321. The other important test is to determine whether a true isothermal halo provides a better model than the pseudo-isothermal halo for the dark matter in late-type disk galaxies, or whether there are better models than either of these. Our analysis of UGC7321 has shown that the vertical gradient of the vertical force provides a significantly stronger constraint on the halo density distribution than does rotation curve decomposition. So, this test can in principle be achieved by analysing UGC7321 and the other galaxies in our sample with both flattened pseudo-isothermal and true isothermal halo models. Such flattened isothermal halos could be flattened by rotation or by anisotropy of the velocity dispersion. This will determine which kind of model is better for both the radial halo force as measured from the rotation curve, and the vertical force of the halo determined from $dK_z/dz$ fitting." }, "1002/1002.1262_arXiv.txt": { "abstract": "{In the 60's, Prof. Str\\\"{o}mgren proposed to use space velocities and ages of moderately young stars to compute their places of formation in order to study the spiral structure in the Galaxy. We have extended this idea to very young stellar clusters in nearby disk galaxies. Near-infrared (NIR) $K$-band images of grand-design spiral galaxies often show bright knots along their spiral arms. Such knots in NGC~2997 have been identified as massive stellar clusters with ages of less than 10~Myr using $JHK$ photometry and $K$-band spectra. Ages of these clusters can be estimated from $JHK$ photometry. Their azimuthal distances from the spiral arms, as measured in the $K$-band, correlates with their ages suggesting that the pattern speed of an underlying density wave can be derived. This method is tested on the grand-design spiral NGC 2997 using VLT data. ", "introduction": "\\citet{stromgren63} suggested to use the migration of young stars with known ages and space velocities to estimate the pattern speed of a density wave in our Galaxy \\citep{lin64, yuan69}. In external galaxies where individual stars cannot be observed, one may consider to use integrated properties of blue, young objects (such as HII regions and OB associations) which often are concentrated in the arms of grand-design spiral galaxies. Strong and very varying attenuation by dust in the arm regions make it difficult to determine reliable, intrinsic colors of such very young objects in visual bands. In the NIR, several spiral galaxies had bright knots along their spiral arms \\citep{gp98}. Such knots in NGC~2997 were identified as very young stellar cluster (ages $<$10~Myr) using $K$-band spectra obtained with ISAAC/VLT \\citep{Grosbol06}. It is possible to estimate accurate ages for clusters younger than $<$8~Myr from their integrated NIR colors and/or \\Brg\\ emission due to the rapid evolution of their high mass stars and the low attenuation by dust. Although space velocities for the clusters cannot be obtained, the general rotation curve of the host galaxy provide enough information to make a crude estimate of their birthplaces assuming that they follow roughly circular orbits. A further advantage of using NIR observations is that phase and shape of a density wave can be measured directly. This opens the way for estimating its pattern speed by comparing ages of individual clusters with their azimuthal offset relative to the spiral perturbation. In addition, one may study star formation induced by such density variations e.g., through large-scale shocks or compressions in the gas \\citep{roberts69}. In the current paper, we study the spatial and color distributions of young stellar clusters in the southern arm of NGC~2997 in order to test the feasibility of this scheme to estimate the pattern speed in external galaxies. \\begin{figure}[t!] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{grosbol_f01.ps}} \\caption{\\footnotesize Direct $K$-band image of the field around the southern arm of NGC~2997 used for the analysis. The scale is indicated by the bar in the lower left corner. } \\label{n2997k} \\end{figure} ", "conclusions": " \\begin{itemize} \\itemsep=0pt \\item the youngest stellar clusters with magnitudes \\MK\\ $<-12$\\,\\mag\\ are concentrated in the arm regions, \\item the bright, young clusters (0.1\\,\\mag\\,$$100 MeV) from \\cyg\\ \\citep{2009Natur.462..620T,2009Sci...326.1512F}. The identification is firm because the detections occur exclusively when \\cyg\\ is flaring in radio and because {\\em Fermi} observations show the HE gamma-ray flux is modulated with the orbital period. The gamma-ray modulation is almost in anti-phase with the X-ray modulation, with the gamma-ray minimum occurring about 0.3-0.4 in phase after X-ray minimum. The modulation amplitude is close to $100\\%$ after background subtraction. The spectrum is consistent with a power law $F_\\nu\\sim \\nu^{-\\alpha}$ with $\\alpha=1.7$. The luminosity above 100 MeV is a few $10^{36} (d/{\\rm 7\\ kpc})^2$ \\eps. Inverse Compton (IC) scattering of photons from the WR star on high energy electrons is a natural candidate to explain the gamma-ray emission. The high temperature of the WR star ($R_\\star\\approx 1$ R$_\\odot$, $T_\\star\\approx 10^5$\\ K) and tight orbit ($d\\approx 3\\ 10^{11}$ cm) imply that the radiation density in photons from the star is $u_{\\star}\\approx 10^5$ erg\\ cm$^{-3}$ at the location of the compact object, which is at least an order-of-magnitude higher than any other X-ray binary. Electrons with Lorentz factors of a few 10$^3$ upscatter 20 eV stellar photons above 100 MeV very efficiently in such a radiation field. IC scattering directly produces a modulation of the flux because of the orbital motion. The maximum occurs when stellar photons are backscattered towards the observer. The accretion disc can also provide seed photons if the HE electrons are close enough. This does not lead to a modulation unless the HE electrons - disk geometry seen by the observer changes with orbital phase \\citep{1977Natur.268..420M}. Pion production is possible if there are high energy protons. However, even in this dense environment, it is less efficient than IC so that its energy requirements are higher. The link between gamma-ray and radio flares suggests that the HE electrons are located in the relativistic jet. Observations of knots in active galactic nuclei show that particles may be accelerated at specific locations along the jet, linked {\\em e.g.} to recollimation shocks \\citep{Stawarz:2006oh}. Assuming the electrons mainly upscatter stellar photons at some location along the jet, the expected IC emission will depend upon the distance to the star, the bulk velocity of the jet and its orientation. This orientation is not necessarily perpendicular to the orbital plane if {\\em e.g.} the inner accretion disc is warped or it depends on the black hole spin axis. However, the jet orientation is fixed as seen by the observer (changing only if the jet precesses). The goal here is to test quantitatively whether the {\\em Fermi} gamma-ray modulation can be reproduced in this framework and to see if constraints can be derived on the jet parameters. % ", "conclusions": "The orbital modulation of the $>$100 MeV flux from \\cyg\\ can be very well fitted by a simple-minded model in which the emission is due to HE electrons up-scattering stellar photons. The HE electrons are situated in two symmetric locations in a relativistic jet with an arbitrary orientation. The fitting procedure reveals that the jet is necessarily inclined to the orbital plane normal. The most likely value is close to the line-of-sight ($\\phi_j\\approx i$, in agreement with the conclusions based on radio imaging of the jet \\citep{2001ApJ...553..766M}. The HE electrons cannot be close to the compact object. They are outside of the system at distances of at least one orbital separation, possibly up to 10$d$. IC scattering of accretion disc photons is then irrelevant. If the compact object in \\cyg\\ is a neutron star, the required power in HE electrons is a significant fraction of the Eddington luminosity. For a black hole, because of the lower system inclination implied, the power required can be as low as $10^{-5} L_{\\rm Edd}$. These conclusions appear robust even when more complex electron distributions and the full IC cross-section are taken into account. Precession can be expected from an inclined jet. It should cause a change in the shape and amplitude of the gamma-ray modulation in the future. The IC cooling timescale is $t_{\\rm ic} \\approx 0.5 (\\gamma_e/10^3)^{-1} (R/d)^2$ seconds (scaled to the orbital separation $d$ and for orbit O1). The size of the gamma ray emitting region is roughly $s\\approx \\beta c t_{\\rm ic}$, giving $s/R\\la 0.04 \\beta(\\gamma_e/10^3)^{-1} (R/d)$ when scaled to $R$. Hence, the assumption that the emission in the {\\em Fermi} energy range is localised holds up to distances $\\approx 10d$ from the star. Cooling slows down at lower energies and electrons emit synchrotron radio beyond the $\\gamma$-ray emission zone on much larger scales. The $\\gamma$-ray emission zone could be related to electron acceleration at a recollimation shock as the jet pushes its way through the stellar wind. The jet is initially over-pressured compared to its environment. It expands freely until its pressure $p_j$ matches that of the environment $p_e$. Here, $p_e$ is the ram pressure of the supersonic wind $\\rho_w v_w^2$. The jet pressure is $p_j\\sim L_j/(\\pi c \\Theta^2 l^2)$ where $L_j$ is the jet power, $\\Theta$ is its opening angle and $l$ is the distance along the jet \\citep[e.g.][]{1997MNRAS.287L...9B}. The pressures equilibrate at \\begin{equation} \\frac{l}{R}\\sim 0.5\\ \\Theta^{-1} {L}_{38}^{1/2} \\dot{M}_{-5}^{-1/2} {v}_{1000}^{-1/2} \\end{equation} with $L_j=10^{38}\\ {\\rm erg\\ s}^{-1}$, $\\dot{M}_{w}=10^{-5} M_\\odot\\ {\\rm yr}^{-1}$ and ${v}_{w}=1000\\ {\\rm km\\ s}^{-1}$. A jet recollimation shock forms beyond $l$. The shock crosses the jet axis after a further distance of order $l$ when the external pressure is constant \\citep{Stawarz:2006oh}. This is roughly the case here since the jet does not extend very far from the system and the dependence of $p_w$ with $l$ remain shallow (unless it is pointed directly away from the star). The location is consistent with the values of $H$ derived above, suggesting this is where jet kinetic or magnetic energy is channeled into particle acceleration. This should be verified by calculations taking into account the non-radial nature of the jet-wind interaction. The shock occurs in the wind only because $\\dot{M}_{w}$ is very large (WR star) and the orbit very tight. Most microquasar jets will actually break out of the immediate vicinity of the system and interact much further away when their pressure matches that of the ISM. Any HE particles there will find a much weaker radiation environment and will be less likely to produce a (modulated) IC gamma-ray flux detectable by {\\em Fermi} or {\\em AGILE}. The emerging picture is that of a jet launched around a black hole, with a moderate bulk relativistic speed, oriented not too far from the line-of-sight, interacting with the WR stellar wind to produce a shock at a distance of 1-10$d$ from the system, where electrons are accelerated to GeV energies and upscatter star photons." }, "1002/1002.0590_arXiv.txt": { "abstract": "We consider the effects of supernovae (SNe) on accretion and star formation in a massive gaseous disk in a large primeval galaxy. The gaseous disk we envisage, roughly 1 kpc in size with $\\gta10^8$M$_\\odot$ of gas, could have formed as a result of galaxy mergers where tidal interactions removed angular momentum from gas at larger radius and thereby concentrated it within the central $\\sim 1$kpc region. We find that SNe lead to accretion in the disk at a rate of roughly 0.1--1 $M_\\odot$ yr$^{-1}$ and induce star formation at a rate of $\\sim 10$--100 M$_\\odot$ per year which contributes to the formation of a bulge; a part of the stellar velocity dispersion is due to SNa shell speed from which stars are formed and a part due to the repeated action of stochastic gravitational field of SNe remnant network on stars. The rate of SNa in the inner kpc is shown to be self regulating, and it cycles through phases of low and high activity. The supernova-assisted accretion transports gas from about one kpc to within a few pc of the center. If this accretion were to continue down to the central black hole then the resulting ratio of BH mass to the stellar mass in the bulge would be of order $\\sim 10^{-2}$--$10^{-3}$, in line with the observed Magorrian relation. ", "introduction": "The CO observations of ultra-luminous infra-red galaxies (ULIGals) find the gas mass in the inner regions of the galaxy to be about 5$\\times10^9$M$_\\odot$, and the average particle density to be of order 10$^3$ cm$^{-3}$ and the kinetic temperature of molecular gas $\\sim 50-100$ K (e.g. Downes \\& Solomon, 1998; see Sanders \\& Mirabel, 1996, for a review). The gas in the central $\\sim$kpc region of the galaxy is likely to have come from distances of the order of 10 kpc when it lost some of its angular momentum due to gravitational tidal torques (Barnes \\& Hernquist, 1992, and references therein; see Barnes, 2002, for a more recent numerical simulation). The inner kiloparsec region of most young massive galaxies is likely composed of a gaseous disk with a mass of several hundred million solar masses, ULIGal being at the extreme end of the mass distribution. This gaseous disk is expected to host star formation at a large rate. Some of these stars will explode and give rise to shock waves in the gaseous disk which will spawn both more star formation and accretion of gas toward the center of the galaxy. We consider these processes analytically in some detail in this paper, paying special attention to the effect they might have on the evolution of the central parts of the galaxy and on the growth of a central black hole. There exists a large body of work on the subject of galaxy mergers, star formation and black hole growth e.g. Sanders et al. (1988), Kauffmann \\& Haehnelt (2000), Kawakatu \\& Umemura (2002), Granato et al. (2004), Croton et al. (2006), Kauffmann \\& Heckman (2009), Chen et al. (2009), (pl. see Kormendy \\& Kennicutt, 2004, for a review), and sophisticated numerical simulations e.g. Barnes \\& Hernquist (1991, 1996), Mihos \\& Hernquist (1996), Di Matteo et al. (2005), Springel et al. (2005), Hopkins et al. (2005), Hopkins \\& Hernquist (2008). What is different in the present work is that we try to capture some of the basic properties of this complex system using analytic results for supernova (SNa) remnant evolution and other simple physical scalings which are hard to capture in numerical simulations due to the large ratio of galaxy size and SNa shell radius. The physical system we consider is described in \\S2 along with the effect SNe have on accretion. Bulge formation as a product of SN-induced star formation in the gaseous disk is discussed in \\S3. The main conclusions and uncertainties of this study can be found in \\S4. ", "conclusions": "We have analyzed the effect of supernovae occurring within the central kpc region of a gaseous disk on the formation of stars and transport of gas from $\\sim 1$kpc to a few pc of the galactic center. The outward transport of angular momentum facilitated by SNe explosions allows for the inward transport of gas that feeds the central black hole. This is a process which may take place quite generically in any galactic disk hosting star formation, although it may not be the dominant process affecting black hole accretion. We have shown that associated with the inward transport of gas swept up by SNe is the shock-induced formation of stars which are born with a random peculiar velocity of $\\sim 10$ km s$^{-1}$. This velocity dispersion increases with time as a result of the stochastic gravitational field associated with filamentary SNa remnants; The stars formed in SNa remnants contribute to the stellar population of a central bulge. Due to the divergent velocity field of an expanding SNa shell there is a maximum length scale for fragmentation of shells or an upper limit to the mass of stars formed in SNa remnants; the SNa-induced stellar IMF is cut-off above a few solar masses and more massive stars can only form if and when a SNa shell slows down to $\\sim10$ km s$^{-1}$. We note that numerous observations have suggested connections, such as we have considered in the present work, between star formation, black hole accretion, and the formation of a stellar bulge. Heckman et al. (2004) suggest that star formation and black hole accretion rates are correlated, and Chen et al. (2009) report an empirical relation between supernova rate and gas accretion rate on the central black hole (see also Xu \\& Wu 2007). Furthermore, Page et al. (2001) report observations suggesting that central black holes and stellar spheroids form concurrently, and Genzel et al. (2006) describe observations of a galaxy hosting both an accreting black hole and a central stellar bulge, with no evidence of a major merger. Hydrodynamic simulations have demonstrated that SNa feedback may produce spherical distributions of stars in dwarf galaxies (Stinson et al. 2009) and in the inner portions of the Galactic bulge (Nakasato \\& Nomoto 2003). We would like to point out the recent work of Wang et al. (2009) that models SNe induced turbulence as an effective viscosity and describes the evolution of a gaseous disk. A limitation of this work is that we have ignored radiative feedback effects which are known to control the steady state accretion rate onto the black hole (e.g. Ostriker et al. 1976, Proga et al 2008, Milosavljevic et al. 2008) and probably also affect the formation rate of stars in the central kpc. Ultimately, large scale simulations resolving the long term evolution of individual SNe remnants, star formation, the multiphase ISM, and the feedback effects of accretion onto a central black hole will be required to more fully elucidate what role SNe-induced accretion and star formation play in galaxy formation." }, "1002/1002.0245_arXiv.txt": { "abstract": "It has already been demonstrated that a mesoscale meteorological model such as Meso-NH (Lafore \\etal~\\cite{laf1998}) is highly reliable in reproducing 3D maps of optical turbulence (Masciadri \\etal~\\cite{mas1999}, Masciadri and Jabouille ~\\cite{mas2001}, Masciadri \\etal~\\cite{mas2004}). Preliminary measurements above the Antarctic Plateau have so far indicated a pretty good value for the seeing: 0.27\" (Lawrence \\etal~\\cite{law2004} ), 0.36\" (Agabi \\etal~\\cite{aga2006}) or 0.4\" (Trinquet \\etal~\\cite{tri2008}) at Dome C. However some uncertainties remain. That's why our group is focusing on a detailed study of the atmospheric flow and turbulence in the internal Antarctic Plateau. Our intention is to use the Meso-NH model to do predictions of the atmospheric flow and the corresponding optical turbulence in the internal plateau. The use of this model has another huge advantage: we have access to informations inside an entire 3D volume which is not the case with observations only. Two different configurations have been used: a low horizontal resolution (with a mesh-size of 100 km) and a high horizontal resolution with the grid-nesting interactive technique (with a mesh-size of 1 km in the innermost domain centered above the area of interest). We present here the turbulence distribution reconstructed by Meso-NH for 16 nights monitored in winter time 2005, looking at the the seeing and the surface layer thickness. ", "introduction": "The extreme low temperatures, the dryness, the typical high altitude of the internal Antarctic Plateau (more than 2500~m), joint to the fact that the optical turbulence seems to be concentrated in a thin surface layer whose thickness is of the order of a few tens of meters do of this site a place in which, potentially, we could achieve astronomical observations otherwise possible only by space. Despite exciting first results (Lawrence \\etal~\\cite{law2004}; Agabi \\etal~\\cite{aga2006}; Trinquet \\etal~\\cite{tri2008}) making the internal Antarctic Plateau a site of potential great interest for astronomical applications, some uncertainties still remain. Here we studied the Dome C area with a mesoscale meteorological model (Meso-NH, Lafore \\etal~\\cite{laf1998}). Numerical simulations offer the advantage to provide volumetric maps of the optical turbulence ($C_N^2$) extended on the whole internal plateau and, ideally, to retrieve comparative estimates in a relative short time and homogeneous way on different places of the plateau. Fifteen winter nights (the same as from Trinquet \\etal~(\\cite{tri2008}) were simulated. Using the forecasted $C_N^2$ profiles, we retrieved the surface layer thicknesses H$_{SL}$ and the free atmosphere seeing ($\\epsilon{_{FA}}$) for all 15 nights. \\par This study is a short survey of the more detailed study available in Lascaux \\etal~(\\cite{las2009}). ", "conclusions": "We studied the performances of the Meso-Nh mesoscale model in reconstructing optical turbulence profiles, looking at the Dome C area, in the internal Antarctic Plateau. This study was focused on the winter season. The results concerning the optical turbulence computations are resolution dependent. A high horizontal resolution mode seems to be mandatory to realize realistic optical turbulence forecast. In high horizontal resolution mode, Meso-Nh gave excellent results: H$_{SL,HIGH}$=48.9$\\pm$ 7.6~m, to be compared to H$_{SL,OBS}$=35.3 $\\pm$ 5.1~m. The resulting free atmosphere seeing is $\\epsilon{_{FA,HIGH}}$=0.35 $\\pm$ 0.24~arcsec, very close to the observed one, $\\epsilon{_{FA,OBS}}$=0.3 $\\pm$ 0.2~arcsec." }, "1002/1002.4300_arXiv.txt": { "abstract": "We use deep HST/ACS observations to calculate the star formation history (SFH) of the Cetus dwarf spheroidal (dSph) galaxy. Our photometry reaches below the oldest main sequence turn-offs, which allows us to estimate the age and duration of the main episode of star formation in Cetus. This is well approximated by a single episode that peaked roughly 12$\\pm$0.5 Gyr ago and lasted no longer than about 1.9$\\pm$0.5 Gyr (FWHM). Our solution also suggests that essentially no stars formed in Cetus during the past 8 Gyrs. This makes Cetus' SFH comparable to that of the oldest Milky Way dSphs. Given the current isolation of Cetus in the outer fringes of the Local Group, this implies that Cetus is a clear outlier in the morphology-Galactocentric distance relation that holds for the majority of Milky Way dwarf satellites. Our results also show that Cetus continued forming stars through $z \\simeq$ 1, long after the Universe was reionized, and that there is no clear signature of the epoch of reionization in Cetus' SFH. We discuss briefly the implications of these results for dwarf galaxy evolution models. Finally, we present a comprehensive account of the data reduction and analysis strategy adopted for all galaxies targeted by the LCID (Local Cosmology from Isolated Dwarfs\\footnotemark[15]) project. We employ two different photometry codes (DAOPHOT/ALLFRAME and DOLPHOT), three different SFH reconstruction codes (IAC-pop/MinnIAC, MATCH, COLE), and two stellar evolution libraries (BaSTI and Padova/Girardi), allowing for a detailed assessment of the modeling and observational uncertainties. \\footnotetext[15]{\\itshape http://www.iac.es/project/LCID} ", "introduction": "\\label{sec:intro} A powerful method to study the mechanisms that drive the evolution of stellar systems is the recovery of their full star formation history (SFH). This can be done by coupling deep and accurate photometry, reaching the oldest main sequence (MS) turn-off (TO), with the synthetic color-magnitude diagram (CMD) modeling technique \\citep{tosi91, bertelli92, tolstoy96, aparicio97, harris01, dolphin02, iacstar, cole07, iacpop}. The details of the first stages of galaxy evolution are particularly interesting, because they directly connect stellar populations research with cosmological studies. For example, the standard paradigm predicts that the role of re-ionization can impact the star formation history of small systems in a measurable way \\citep[e.g.,][]{i86, r86, e92, br92, cn94, qke96, tw96, kbs97, barkana99, bullock00, tassis03, ricotti05, gnedin06, okamoto09}. In this context, we designed a project aimed at recovering the full SFHs of six isolated galaxies in the Local Group (LG), namely IC~1613, Leo~A, LGS~3, Phoenix, Cetus, and Tucana, with particular focus on the details of the early SFH. Isolated systems were selected because they are thought to have completed few, if any, orbits inside the LG. Therefore, compared to satellite dSphs, they spent most of their lifetime unperturbed, so that their evolution is expected not to be strongly affected by giant galaxies. A general summary of the goals, design and outcome of the LCID project can be found in Gallart et al. (in prep), where a comparative study of the results for the six galaxies is presented. In this paper we focus on the stellar populations of the Cetus dSph galaxy only. This galaxy was discovered by \\citet{whiting99}, by visual inspection of the ESO/SRC plates. From follow-up CCD observations they obtained a CMD to a depth of $V \\sim 23$, clearly showing the upper part of the red giant branch (RGB), but no evidence of a young main sequence, nor of either red or blue supergiants typical of evolved young populations. From the tip of the RGB (TRGB) they could estimate a distance modulus $(m-M)_0 = 24.45 \\pm 0.15$ mag ($776 \\pm 53$ kpc), assuming $E(B-V) = 0.03$ from the \\citet{schlegel98} maps and an absolute magnitude $M_{I}(TRGB) = -3.98 \\pm 0.05$ mag for the tip. The color of the RGB stars near the tip allowed a mean metallicity estimate of $[Fe/H] = -1.7$ dex with a spread of $\\sim 0.2$ dex. Deeper HST/WFPC2 data presented by \\citet{sarajedini02} confirmed the distance modulus of Cetus: $(m-M)_0 = 24.49 \\pm 0.14$ mag, or $790 \\pm 50$ kpc, assuming $M_{I}(TRGB) = -4.05 \\pm 0.10$ mag. Based on the color distribution of bright RGB stars, they estimated a mean metallicity $[Fe/H] = -1.9$ dex. An interesting feature of the Cetus CMD presented by \\citet{sarajedini02} is that the horizontal branch (HB) seems to be more populated in the red than in the blue part \\citep[see also the CMD by][based on VLT/FORS1 data]{tolstoycetus00}. \\citet{sarajedini02} calculated an HB $index = -0.91 \\pm 0.09$\\footnotemark[16]. By using an empirical relation between the mean color of the HB and its morphological type, derived from globular clusters of metallicity similar to that of Cetus, they concluded that the Cetus HB is too red for its metallicity, and therefore, is to some extent affected by the so-called second parameter. Assuming that this is mostly due to age, they proposed that Cetus could be 2-3 Gyr younger than the old Galactic globular clusters. \\footnotetext[16]{The HB index is defined as ($B$-$R$)/($B$+$V$+$R$), \\citep{lee90}, where B and R are the number of HB stars bluer and redder than the instability strip, and V is the total number of RR Lyrae stars.} Recent work by \\citet{lewis07} presented wide-field photometry from the Isaac Newton Telescope together with Keck spectroscopic data for the brightest RGB stars. From CaT measurements of 70 stars, they estimated a mean metallicity of $[Fe/H] = -1.9$ dex, fully consistent with previous estimates based on photometric data. Moreover, they determined for the first time a systemic heliocentric velocity of $\\sim$ 87 km~s$^{-1}$, an internal velocity dispersion of 17 km~s$^{-1}$ and, interestingly enough, little evidence of rotation. The search for H~I gas \\citep{bouchard06} revelead three possible candidate clouds associated with Cetus, located beyond its tidal radius, but within 20 Kpc from its centre. Considering the velocity of these clouds, \\citet{bouchard06} suggested that the velocity of Cetus should be of the order of -280$\\pm 40 km~s^{-1}$ to have a high chance of association. Clearly, this does not agree with the estimates later given by \\citet{lewis07}, who pointed out that Cetus is devoid of gas, at least to the actual observational limits. Cetus is particularly interesting for two distinctive characteristics. First, the high degree of isolation, it being at least $\\sim 680$ kpc away from both the Galaxy and M31. Cetus is, together with Tucana, one of the two isolated dwarf spheroidals in the LG. Moreover, based on wide-field INT telescope data, \\citet{mcconnachie06} estimated a tidal radius of 6.6 kpc, which would make Cetus the largest dSph in the LG. Therefore, an in-depth study of its stellar populations, and how they evolved with time, could provide important clues about the processes governing dwarf galaxy evolution. The paper is organized as follows. In \\S \\ref{sec:obse} we present the data and the reduction strategy. \\S \\ref{sec:cmd} discusses the details of the derived color-magnitude diagram (CMD). In \\S \\ref{sec:methods} we summarize the methods adopted to recover the SFH, and we discuss the effect of the photometry/library/code on the derived SFH. The Cetus SFH is presented in \\S \\ref{sec:results}, where we discuss the details of the results, as well as the tests we performed to estimate the duration of the main peak of star formation and the radial gradients. A discussion of these results in the general context of galaxy evolution are presented in \\S \\ref{sec:discu}, and our conclusions are summarized in \\S \\ref{sec:summa}. ", "conclusions": "" }, "1002/1002.1713_arXiv.txt": { "abstract": "{} {We outline the Bayesian approach to inferring \\fnl, the level of \\ngs\\ of local type. Phrasing \\fnl\\ inference in a Bayesian framework takes advantage of existing techniques to account for instrumental effects and foreground contamination in CMB data and takes into account uncertainties in the cosmological parameters in an unambiguous way.} {We derive closed form expressions for the joint posterior of \\fnl\\ and the reconstructed underlying curvature perturbation, $\\Phi$, and deduce the conditional probability densities for \\fnl\\ and $\\Phi$. Completing the inference problem amounts to finding the marginal density for \\fnl. For realistic data sets the necessary integrations are intractable. We propose an exact Hamiltonian sampling algorithm to generate correlated samples from the \\fnl\\ posterior. For sufficiently high signal-to-noise ratios, we can exploit the assumption of weak \\ngy\\ to find a direct Monte Carlo technique to generate \\emph{independent} samples from the posterior distribution for \\fnl. We illustrate our approach using a simplified toy model of CMB data for the simple case of a 1-D sky.} {When applied to our toy problem, we find that, in the limit of high signal-to-noise, the sampling efficiency of the approximate algorithm outperforms that of Hamiltonian sampling by two orders of magnitude. When \\fnl\\ is not significantly constrained by the data, the more efficient, approximate algorithm biases the posterior density towards $f_{\\NL} = 0$.} {} ", "introduction": "\\label{sec:intro} The analysis of cosmic microwave background (CMB) radiation data has considerably improved our understanding of cosmology and played a crucial role in constraining the set of fundamental cosmological parameters of the universe \\citep{2007ApJS..170..377S, 2009ApJS..180..225H}. This success is based on the intimate link between the temperature fluctuations we observe today and the physical processes taking place in the very early universe. Inflation is currently the favored theory predicting the shape of primordial perturbations \\citep{1981PhRvD..23..347G, 1982PhLB..108..389L}, which in its canonical form leads to very small \\ngs\\ that are far from being detectable by means of present-day experiments \\citep{Maldacena:2002vr, Acquaviva:2002ud}. However, inflation scenarios producing larger amounts of \\ngy\\ can naturally be constructed by breaking one or more of the following properties of canonical inflation: slow-roll, single-field, Bunch-Davies vacuum, or a canonical kinetic term \\citep{2004PhR...402..103B}. Thus, a positive detection of primordial \\ngy\\ would allow us to rule out the simplest models. Combined with improving constraints on the scalar spectral index $n_s$, the test for \\ngy\\ is therefore complementary to the search for gravitational waves as a means to test the physics of the early Universe. A common strategy for estimating primordial \\ngy\\ is to examine a cubic combinations of filtered CMB sky maps \\citep{2005ApJ...634...14K}. This approach takes advantage of the specific bispectrum signatures produced by primordial \\ngy\\ and yields to a computationally efficient algorithm. When combined with the variance reduction technique first described by \\citet{2006JCAP...05..004C} these bispectrum-based techniques are close to optimal, where optimality is defined as saturation of the Cramer-Rao bound. Lately, a more computationally costly minimum variance estimator has been implemented and applied to the WMAP5 data \\citep{2009JCAP...09..006S}. Recently, a Bayesian approach has been introduced in CMB power spectrum analysis and applied successfully to WMAP data making use of Gibbs-sampling techniques \\citep{2004ApJ...609....1J, 2004PhRvD..70h3511W}. Within this framework, one draws samples from the posterior probability density given the data without explicitly calculating it. The target probability distribution is finally constructed out of the samples directly, thus computationally costly evaluations of the likelihood function or its derivatives are not necessary. Another advantage of the Bayesian analysis is that the method naturally offers the possibility to include a consistent treatment of the uncertainties associated with foreground emission or instrumental effects \\citep{2008ApJ...672L..87E}. As it is possible to model CMB and foregrounds jointly, statistical interdependencies can be directly factored into the calculations. This is not straightforward in the frequentist approach where the data analysis is usually performed in consecutive steps. Yet another important and desirable feature is the fact that a Bayesian analysis obviates the necessity to specify fiducial parameters, whereas in the frequentist approach it is only possible to test one individual null hypothesis at a time. In this paper, we pursue the modest goal of developing the formalism for the extension of the Bayesian approach to the analysis of \\ngn\\ signals, in particular to local models, where the primordial perturbations can be modeled as a spatially local, non-linear transformation of a \\gn\\ random field. Utilizing this method, we are able to write down the full posterior probability density function (PDF) of the level of \\ngy. We demonstrate the principal aspects of our approach using a 1-D toy sky model. Although we draw our discussion on the example of CMB data analysis, the formalism presented here is of general validity and may also be applied within a different context. The paper is organized as follows. In \\sect{sec:basics} we give a short overview of the theoretical background used to characterize primordial perturbations. We present a new approximative approach to extract the amplitude of \\ngs\\ from a map in \\sect{sec:sampling} and verify the method by means of a simple synthetic data model (\\sect{sec:model}). We compare the performance of our technique to an exact Hamiltonian Monte Carlo sampler which we developer in \\sect{sec:hmc} and discuss the extensions of the model required to deal with a realistic CMB sky map (\\sect{sec:cmb}). Finally, we summarize our results in \\sect{sec:summary}. ", "conclusions": "\\label{sec:model} To verify our results and demonstrate the applicability of the method, we implemented a simple 1-D toy model. We considered a vector $\\Phi_{\\LL}$ of random numbers generated from a heptadiagonal covariance matrix with elements \\begin{equation} P_\\Phi = \\begin{pmatrix} & & & \\ddots \\\\ \\dots \\, 0 & 0.1 & 0.2 & 0.5 & 1.0 & 0.5 & 0.2 & 0.1 & 0 \\, \\dots \\\\ & & & & & \\ddots \\\\ \\end{pmatrix} \\times 10^{-10} \\, . \\end{equation} Then, a data vector with weak \\ngy\\ according to \\eq{eq:fnldef} was produced and superimposed with \\gn\\ white noise. Constructed in this way, it is of the order $\\mathcal{O}(10^{-5})$, thus the amplitude of the resulting signal $s$ is comparable to CMB anisotropies. The data vector had a length of $10^6$ pixels; for simplicity, we set the beam function $B$ and the linear transformation matrix $M$ to unity. This setup allows a brute force implementation of all equations at a sufficient computational speed. We define the signal-to-noise ratio ($S/N$) per pixel as the standard deviation of the input signal divided by the standard deviation of the additive noise. It was chosen in the range 0.5-10 to model the typical S/N per pixel of most CMB experiments. To reconstruct the signal, we draw $1000$ samples according to the scheme in \\eq{eq:scheme}. Whereas the $\\Phi_{\\NL}$ can be generated directly from a simple \\gn\\ distribution with known mean and variance, the construction of the \\fnl\\ is slightly more complex. For each $\\Phi_{\\NL}$, we ran a Metropolis Hastings algorithm with symmetric \\gn\\ proposal density with a width comparable to that of the target density and started the chain at $f_{\\NL} = 0$. We run the \\fnl\\ chain to convergence. We ensured that after ten accepted steps the sampler has decorrelated from the starting point. Our tests conducted with several chains run in parallel give $1 < R < 1.01$, where R is the convergence statistic proposed by \\citet{199211}. We record the last element of the chain as the new \\fnl\\ sample. Finally, we compared the obtained sets of values $\\{\\Phi_{\\NL}^i\\}$, $\\{f_{\\NL}^i\\}$ to the initial data. An example is shown in \\fig{fig:ex_phi_l}, where we illustrate the reconstruction of a given potential $\\Phi_{\\NL}$ for different signal-to-noise ratios per pixel. The $1-\\sigma$ error bounds are calculated from the 16~\\% and 84~\\% quantile of the generated sample. Typical posterior densities for \\fnl\\ as derived from the sample can be seen in \\fig{fig:ex_f_nl_1}. We considered the cases $f_{\\NL} = 0$ and $f_{\\NL} = 200$ with $S/N = 10$ per pixel and show the distributions generated from $1000$ draws. The derived posterior densities possesses a mean value of $f_{\\NL} = 6 \\pm 40$ and $f_{\\NL} = 201 \\pm 40$, respectively. The width of the posterior is determined by both the shape of the conditional PDF of \\fnl\\ for a given $\\Phi_{\\NL}$ and the shift of this distribution for different draws of $\\Phi_{\\NL}$ (\\fig{fig:ex_f_nl_2}). The analysis of several data sets indicate that the approximation does not bias the posterior density if the data are decisive. We illustrate this issue in the left panel of \\fig{fig:f_nl_mean_error}, where we show the distribution of the mean values $\\langle f_{\\NL} \\rangle$ of the posterior density constructed from 100 independent simulations. For an input value of $f_{\\NL} = 200$ we derive a mean value $\\langle f_{\\NL} \\rangle = 199.3 \\pm 34.8$ and conclude that our sampler is unbiased for these input parameters. For a high noise level, however, the $\\Phi_{\\NL}$ can always be sampled such that they are purely \\gn\\ fields and thus the resulting PDF for \\fnl\\ is then shifted towards $f_{\\NL} = 0$. This behavior is demonstrated in \\fig{fig:f_nl_low_sn} where we compare the constructed posterior density for the cases $S/N = 10$ and $S/N = 0.5$ per pixel. If the noise level becomes high, the approximated prior distribution dominates and leads to both, a systematic displacement and an artificially reduced width of the posterior. Therefore, the sampler constructed here is conservative in a sense that it will tend to underpredict the value of \\fnl\\ if the data are ambiguous. An example of the evolution of the drawn \\fnl\\ samples with time can be seen in \\fig{fig:ex_f_nl_3}, where we in addition show the corresponding autocorrelation function as defined via \\begin{equation} \\xi (\\Delta N) = \\frac{1}{N} \\sum^{N}_{i} \\frac{(f_{\\NL}^i - \\mu) \\cdot (f_{\\NL}^{i + \\Delta N} - \\mu)}{\\sigma^2} \\, , \\end{equation} where N is the length of the generated \\fnl\\ chain with mean $\\mu$ and variance $\\sigma^2$. The uncorrelated samples of \\fnl\\ ensure an excellent mixing of the chain resulting in a fast convergence rate. \\label{sec:summary} In this paper, we developed two methods to infer the amplitude of the \\ngy\\ parameter \\fnl\\ from a data set within a Bayesian approach. We focused on the so called local type of \\ngy\\ and derived an expression for the joint probability distribution of \\fnl\\ and the primordial curvature perturbations, $\\Phi$. Despite the methods are of general validity, we tailored our discussion to the case example of CMB data analysis. We developed an exact Markov Chain sampler that generates correlated samples from the joint density using the Hamiltonian Monte Carlo approach. We implemented the HMC sampler and applied it to a toy model consisting of simulated measurements of a 1-D sky. These simulations demonstrate that the recovered posterior distribution is consistent with the level of simulated \\ngy. With two approximations that exploit the fact that the \\ngn\\ contribution to the signal is next order in perturbation theory, we find a far more computationally efficient Monte Carlo sampling algorithm that produces \\emph{independent} samples from the \\fnl\\ posterior. The regime of applicability for this approximation is for data with high signal-to-noise and weak \\ngy. By comparison to the exact HMC sampler, we show that our approximate algorithm reproduces the posterior location and shape in its regime of applicability. If non-zero \\fnl\\ is not supported by the data the method is biased towards Gaussianity. The approximate posterior more strongly prefers zero \\fnl\\ compared to non-zero values than the exact posterior, as expected given the nature of the approximations which Gaussianize the prior. This method is therefore only applicable if the data contains sufficient support for the presence of \\ngy\\, essentially overruling the preference for Gaussianity in our approximate prior. Our efficient method enables us to perform a Monte Carlo study of the behavior of the posterior density for our toy model data with high signal-to-noise per pixel. We found that the width of the posterior distribution does not change as a function of the level of \\ngy\\ in the data, contrary to the frequentist estimator where there is an additional, \\fnl\\ dependent, variance component \\citep{2007JCAP...03..019C, 2007PhRvD..76j5016L}. Our results suggest that this may be an advantage of the Bayesian approach compared to the frequentist approach, motivating further study of the application of Bayesian statistics to the search for primordial local \\ngy\\ in current and future CMB data. We close on a somewhat philosophical remark. Even though we chose a Gaussian prior approximation for expediency, it may actually be an accurate model of prior belief for many cosmologists since canonical theoretical models predict Gaussian perturbations. From that perspective our fast, approximate method may offer some (philosophically interesting) insight into the question ``what level of signal-to-noise in the data is required to convince someone of the presence of \\ngy\\ whose prior belief is that the primordial perturbations are Gaussian?''" }, "1002/1002.1239_arXiv.txt": { "abstract": "{Supersonic turbulence in molecular clouds is a key agent in generating density enhancements that may subsequently go on to form stars. The stronger the turbulence -- the higher the Mach number -- the more extreme the density fluctuations are expected to be. Numerical models predict an increase in density variance, $\\sigma^{2}_{\\rho/\\rho_{0}}$, with rms Mach number, $M$ of the form: $\\sigma^{2}_{\\rho/\\rho_{0}} = b^{2}M^{2}$, where $b$ is a numerically-estimated parameter, and this prediction forms the basis of a large number of analytic models of star formation. We provide an estimate of the parameter $b$ from $^{13}$CO J=1--0 spectral line imaging observations and extinction mapping of the Taurus molecular cloud, using a recently developed technique that needs information contained solely in the projected column density field to calculate $\\sigma^{2}_{\\rho/\\rho_{0}}$. When this is combined with a measurement of the rms Mach number, $M$, we are able to estimate $b$. We find $b = 0.48^{+0.15}_{-0.11}$, which is consistent with typical numerical estimates, and is characteristic of turbulent driving that includes a mixture of solenoidal and compressive modes. More conservatively, we constrain $b$ to lie in the range 0.3--0.8, depending on the influence of sub-resolution structure and the role of diffuse atomic material in the column density budget (accounting for sub-resolution variance results in higher values of $b$, while inclusion of more low column density material results in lower values of $b$; the value $b = 0.48$ applies to material which is predominantly molecular, with no correction for sub-resolution variance). We also report a break in the Taurus column density power spectrum at a scale of $\\sim$~1~pc, and find that the break is associated with anisotropy in the power spectrum. The break is observed in both $^{13}$CO and dust extinction power spectra, which, remarkably, are effectively identical despite detailed spatial differences between the $^{13}$CO and dust extinction maps. } ", "introduction": "Recent years have seen a proliferation of analytical models of star formation that provide prescriptions for star formation rates and initial mass functions based on physical properties of molecular clouds (Padoan \\& Nordlund 2002; Krumholz \\& McKee 2005; Elmegreen 2008; Hennebelle \\& Chabrier 2008; Padoan \\& Nordlund 2009; Hennebelle \\& Chabrier 2009). While these models differ in their details, they are all fundamentally based on the same increasingly influential idea that has emerged from numerical models: that the density PDF is lognormal in form for isothermal gas (V\\'{a}zquez-Semadeni 1994), with the normalized density variance increasing with the rms Mach number ($\\sigma^{2}_{\\rho/\\rho_{0}} = b^{2}M^{2}$ where $\\rho$ is the density, $\\rho_{0}$ is the mean density, M is the 3D rms Mach number, and $b$ is a numerically determined parameter; Padoan, Nordlund, \\& Jones 1997). There is some uncertainty on the value of $b$. Padoan et al (1997) propose $b = 0.5$ in 3D, while Passot \\& V\\'{a}zquez-Semadeni (1998) found $b = 1$ using 1D simulations. Federrath, Klessen, \\& Schmidt (2008) suggested that $b = 1/3$ for solenoidal (divergence-free) forcing and $b = 1$ for compressive (curl-free) forcing in 3D. A value of $b = 0.25$ was recently found by Kritsuk et al (2007) in numerical simulations that employed a mixture of compressive and solenoidal forcing. Lemaster \\& Stone (2008) found an non-linear relation between $\\sigma^{2}_{\\rho/\\rho_{0}}$ and $M^{2}$, but it is very similar to a linear relation with $b = 1/3$ over the range of Mach numbers they analyzed. They also found that magnetic fields appear to have only a weak effect on the $\\sigma^{2}_{\\rho/\\rho_{0}}$~--~$M^{2}$ relation, and noted that the relation may be different in conditions of decaying turbulence. Currently, observational information on the value of $b$, or the linearity of the $\\sigma^{2}_{\\rho/\\rho_{0}}$~--~$M^{2}$ relation, is very sparse. A major obstacle is the inaccessibility of the 3D density field. Recently, Brunt, Federrath, \\& Price (2009; BFP) have developed and tested a method of calculating $\\sigma^{2}_{\\rho/\\rho_{0}}$ from information contained solely in the projected column density field, which we apply to the Taurus molecular cloud in this paper. Using a technique similar to that of BFP, Padoan, Jones, \\& Nordlund (1997) have previously estimated a value of $b = 0.5$ ($M = 10$, $\\sigma_{\\rho/\\rho_{0}} = 5$) for the IC~5146 molecular cloud. In sub-regions of the Perseus molecular cloud Goodman, Pineda, \\& Schnee (2009) found no obvious relation between Mach number and normalized column density variance, $\\sigma^{2}_{N/N_{0}}$. It has been suggested by Federrath et al (2008) that this could be due to differing levels of compressive forcing, but it may also be due to differing proportions of $\\sigma^{2}_{\\rho/\\rho_{0}}$ projected into $\\sigma^{2}_{N/N_{0}}$ (BFP). Alternatively, this may imply that there is no obvious relation between $\\sigma^{2}_{\\rho/\\rho_{0}}$ and $M^{2}$. The aim of this paper is to constrain the $\\sigma^{2}_{\\rho/\\rho_{0}}$~--~$M^{2}$ relation observationally. In Section~2, we briefly describe the BFP technique, and in Section~3 we apply it to spectral line imaging observations (Narayanan et al 2008; Goldsmith et al 2008) and dust extinction mapping (Froebrich et al 2007) of the Taurus molecular cloud to establish an observational estimate of $b$. Section~4 provides a summary. ", "conclusions": "We have provided an estimate of the parameter $b$ in the proposed $\\sigma^{2}_{\\rho/\\rho_{0}} = b^{2}M^{2}$ relation for supersonic turbulence in isothermal gas, using $^{13}$CO observations of the Taurus molecular cloud. Using $^{13}$CO and 2MASS extinction data, we find $b = 0.48^{+0.15}_{-0.11}$, which is consistent with the value originally proposed by Padoan et al (1997) and characteristic of turbulent forcing which includes a mixture of both solenoidal and compressive modes (Federrath et al 2009). Our value of $b$ is a lower limit if significant structure exists below the resolution of our observations (effectively, below the sonic scale). If diffuse material is included in the column density budget then we find a somewhat lower value of $b = 0.32$, although in this case there are some questions over the assumption of isothermality in the Mach number calculation. Our most conservative statement is therefore that $b$ is constrained to lie in the range $0.3 \\lesssim b \\lesssim 0.8$, which is comparable to the range of current numerical estimates. However, the $\\sigma^{2}_{\\rho/\\rho_{0}} - M^{2}$ relation has been tested to only relatively low Mach numbers ($M \\lesssim 7$) in comparison to that in Taurus, so further numerical exploration of the high Mach number regime should be carried out. In principle, gravitational amplification of turbulently-generated density enhancements could increase $\\sigma^{2}_{\\rho/\\rho_{0}}$ at roughly constant Mach number, with gravity possibly acting analogously to compressive forcing. Numerical investigation of the role of gravity in the $\\sigma^{2}_{\\rho/\\rho_{0}} - M^{2}$ would be worthwhile. Further analysis of a larger sample of molecular clouds is needed to investigate the linearity of the $\\sigma^{2}_{\\rho/\\rho_{0}}$~--~$M^{2}$ relation and possible variations in $b$. As our main sources of uncertainty arise because of questions over how to treat the diffuse (likely atomic) regime and how to account for unresolved density variance, further progress on improving the accuracy of measuring $b$ ultimately must involve answering what is meant by ``a cloud'' (Ballesteros-Paredes, V\\'{a}zquez-Semadeni, \\& Scalo 1999) and down to what spatial scale is significant structure likely to be present?" }, "1002/1002.1525_arXiv.txt": { "abstract": " ", "introduction": "\\label{sec:PN_intro} A classical planetary nebula (PN) is a shell of ionized gas, ejected from a low- to intermediate-mass star ($\\sim$1 to 8 M$_{\\odot}$) towards the end of its life. A PN is produced at the end of the asymptotic giant branch (AGB) phase, as the red giant ejects its outer envelope in a final stage of copious mass loss termed the `superwind'. After the envelope ejection, the remnant core of the star increases in temperature before the nuclear burning ceases and the star quickly fades, becoming a white dwarf (WD). The high temperature of the central star (CS) causes the previously ejected material to be ionized, which becomes visible as a PN, its shape sculpted by the interaction between the old red giant envelope and a tenous, fast wind from the hot CS (Kwok, Purton \\& Fitzgerald 1978). PN\\footnote{For simplicity, we use PN as an abbreviation for both singular and plural forms.} are an important, albeit brief ($\\leq$10$^5$ yr), evolutionary phase in the lifetimes of a significant fraction of Milky Way stars. They are an important tool in our understanding of the physics of mass loss for intermediate-mass stars, the chemical enrichment of our Galaxy, and in turn, its star formation history. PN are ideal test particles to probe the dynamics of the Milky Way and are amongst the best kinematic tracers in external galaxies. The PN luminosity function is also a powerful extragalactic distance indicator (e.g. Ciardullo 2010, this issue). The total number of currently known Galactic PN is nearly 3000, which includes all known {\\sl confirmed} PN from the catalogues of Acker et al. (1992, 1996) and Kohoutek (2001), published PN found since 2000, including $\\sim$1250 from the recent MASH catalogues (Parker et al. 2006; Miszalski et al. 2008), additional discoveries from the IPHAS survey (Viironen et al. 2009a, 2009b), the new faint PN found as part of the ongoing Deep Sky Hunters project (Kronberger et al. 2006; Jacoby et al. 2010, this issue), plus a modest number of additional PN gleaned from the literature. It is beyond the scope of this paper to give an exhaustive summary of our current state of knowledge concerning PN. For reviews on different aspects of the PN phenomenon, the reader is referred to the works of Pottasch (1992), Balick \\& Frank (2002), and De Marco (2009), and the monograph by Kwok (2000). ", "conclusions": "The total number of true, likely and possible PN now known in the Milky Way is nearly 3000, almost double the number known a decade ago, and with the potential to grow more over the next few years. These discoveries are a legacy of the recent availability of wide field, narrowband, imaging surveys. Many types of objects can mimic PN, and consequently our current Galactic and extragalactic samples are liberally salted with emission-line objects and other sources masquerading as PN. Many contaminants still remain to be excised as PN, particularly from the pre-MASH catalogues. On the plus side, we find that much improved discrimination of true PN from their mimics is now possible based on the wide variety of high-quality UV, optical, IR, and radio data that is now available. We recommend several diagnostic plots for classification, but it pays to be aware of their limitations, as it does to be cognisant of the full zoo of potential mimics. While data-mining has its place, we caution that unusual objects be classified on a case-by-case basis. On the downside is the realisation that the overall `PN phenomenon' probably includes a melange of subgroups that observationally have subtle differences, though we think these are quite rare. It appears likely that these subgroups have resulted from a range of discrete evolutionary channels. In other words, the PN phenomenon appears to be a mixed bag. In particular, the role of binarity in the formation and shaping of PN is still an open question (Soker \\& Livio 1989; Soker 1997, 1998; Moe \\& De Marco 2006; De Marco 2009; Miszalski et al. 2009b), but it appears likely that many CS appear to be single stars (Soker 2002; Soker \\& Subag 2005). We reiterate that there is a surprising diversity in the spectra of emission-line CS of PN-like nebulae, but the exact relationships between these nuclei and \\emph{some} symbiotic and B[e] stars is still unclear. Future observational and theoretical advances will allow a more precise taxonomy of the PN phenomenon, but a consensus eludes the community at this point in time." }, "1002/1002.3650_arXiv.txt": { "abstract": "We present a new, magnetohydrodynamic mechanism for inflation of close-in giant extrasolar planets. The idea behind the mechanism is that current, which is induced through interaction of atmospheric winds and the planetary magnetic field, results in significant Ohmic dissipation of energy in the interior. We develop an analytical model for computation of interior Ohmic dissipation, with a simplified treatment of the atmosphere. We apply our model to HD209458b, Tres-4b and HD189733b. With conservative assumptions for wind speed and field strength, our model predicts a generated power that appears to be large enough to maintain the transit radii, opening an unexplored avenue towards solving a decade-old puzzle of extrasolar gas giant radius anomalies. ", "introduction": "The detection of the first transiting extrasolar planet HD209458b \\citep{2000ApJ...529L..45C,2000ApJ...529L..41H} marked the first observation of a planet whose radius is anomalously large. With the current aggregate of transiting planets exceeding 60, over-inflated ``hot Jupiters\" are now known to be common (Fig.1), and understanding their radii has become recognized as an outstanding problem in planetary astrophysics \\citep{2010RPPh...73a6901B}. Most proposed explanations require an interior power source that would replace the radiated heat from gravitational contraction and cause a planet to reach thermal equilibrium with a larger-than-expected radius. In the context of such solutions, the generated heat must be deposited into the interior envelope, i.e. below the radiative/convective boundary, in order to maintain the core entropy (and therefore the radius) of the planet. Notably, eccentricity tides \\citep{2001ApJ...548..466B}, obliquity tides of a Cassini state \\citep{2005ApJ...628L.159W}, and deposition of kinetic energy to adiabatic depths by dynamical and convective instabilities \\citep{2002A&A...385..156G} have been invoked to provide an extra power source in the interior of the planet. It has been shown that the required powers are rather modest \\citep{2007ApJ...661..502B}, but it is unlikely that any of the proposed solutions alone are able to account for all observed radii \\citep{2010RPPh...73a6901B,2009SSRv..tmp..108F}. Here we show that the anomalous sizes of close-in exo-planets can be explained by a magnetohydrodynamic mechanism. The interactions of zonal winds with the expected planetary magnetic field in a thermally ionized atmosphere induce an emf that drives electrical currents into the interior. These currents dissipate Ohmically and thus maintain the interior entropy of the planet. The primary controlling factors in our model are the atmospheric temperature, wind velocity and strength of the magnetic field, as they dictate how much current is allowed to penetrate the interior. Other variables, such as metalicity also contribute, but to a smaller degree. Our results predict that interior heating of this kind occurs in all close-in exoplanets with magnetic fields, but it is negligible if the atmospheric temperature is not high enough for sufficient thermal ionization to take place. Smaller, but hot exoplanets are attributed to heavy element enrichment in the interior. While the inflation mechanism we present here is general, the quantitative modeling in this work is specific to HD209458b, Tres-4b, and HD189733b which are arguably the better studied transiting exoplanets. \\begin{figure} \\epsscale{1.0} \\plotone{f1.pdf} \\caption{Scatter-plot of mass vs. radius of transiting Jovian exo-planets. The three planets considered in the text as well as Jupiter \\& Saturn are labeled. The two lines represent the theoretical mass-radius relationships for a core-less planet (dashed) and one with a $40M_{\\oplus}$ core (solid) from \\cite{2003ApJ...592..555B}. Planets above the dashed line require an inflation mechanism to halt gravitational contraction.} \\end{figure} ", "conclusions": "In this letter, we have presented a new, magnetohydrodynamic mechanism for inflation of extrasolar gas giants. Our calculations show that the heating, necessary to maintain the seemingly anomalous radii of transiting exo-planets, naturally emerges from considerations of interactions between partially ionized winds and the planetary magnetic field. Interestingly, there seems to be a set limit to the extent that Ohmic dissipation can heat the interior, making this theory testable. Currently, there is significant uncertainty with respect to the calculation of the required interior heating, because core masses are unknown. However, dynamical determinations of interior structure \\citep{2009ApJ...704L..49B,2009ApJ...698.1778R} may allow us to resolve the degeneracy for a fraction of observed planets, and provide a solid test-bed for the mechanism we've presented here. There is a number of interesting additional questions that our model inevitably brings up. First, recall that our treatment of the induction equation is kinematic. In reality, flow modification by the Lorenz force may play an important role in determining the actual wind patterns. While this effect may be small for HD209458b and HD189733b, weather on hotter planets, such as Tres-4b or Wasp-12b may be more intimately linked with their magnetic fields, calling for a magentohydrodynamic treatment of the atmospheric circulation. Generally, when zonal winds interact with a background dipole field, they give rise to poloidal current which in turn, gives rise to a predominantly toroidal, unobservable field. However, the dayside-to-nightside flows that are present at higher levels in the atmosphere may modify the flow in an interesting way that may eventually be astronomically observable. Second, we are neglecting the stellar magnetic field. The star's magnetic field is likely to be considerably smaller than the planetary field at the planetary orbital radius, but induction by stellar field as well as coupling of the stellar and planetary magnetic field lines is certainly plausible. This too, may produce an astronomically observable signature. Finally, we are neglecting the effects the induced current in the interior on the planetary dynamo. Considerations of this sort may influence the background magnetic field of the planet. All of these aspects call for a self-consistent treatment of the full problem. Such calculations would no-doubt provide further insight into the physical structure of extra-solar gas giants. \\begin{appendix} \\begin{table} \\begin{center} \\caption{Ohmic dissipation acquired at various pressures in various models of HD209458b, Tres-4b, and HD189733b} \\begin{tabular}{llllllll} \\tableline\\tableline Planet & Y & $T_{iso}$ (K) & Z ($\\times$solar) & $\\mathbb{P}$ [$P<10$ Bars] (W) & $\\mathbb{P}$ [$P>10$ Bars] (W) & $\\mathbb{P}$ [$P>100$ Bars] (W) \\\\ \\tableline HD209458b & 0.24 & 1400 & 1 & $2.30\\times10^{19}$ & $2.23\\times10^{17}$ & $1.09\\times10^{16}$ \\\\ HD209458b & 0.24 & 1400 & 10 & $7.28\\times10^{19}$ & $7.06\\times10^{17}$ & $3.43\\times10^{16}$ \\\\ HD209458b & 0.24 & 1700 & 1 & $1.14\\times10^{21}$ & $1.01\\times10^{19}$ & $5.60\\times10^{17}$ \\\\ HD209458b & 0.24 & 1700 & 10 & $3.61\\times10^{21}$ & $3.19\\times10^{19}$ & $1.77\\times10^{18}$ \\\\ HD209458b & 0.24 & 2000 & 1 & $1.22\\times10^{22}$ & $3.24\\times10^{20}$ & $7.09\\times10^{19}$ \\\\ HD209458b & 0.24 & 2000 & 10 & $3.89\\times10^{22}$ & $1.05\\times10^{21}$ & $2.29\\times10^{20}$ \\\\ HD209458b & 0.3 & 1400 & 1 & $2.22\\times10^{19}$ & $1.30\\times10^{17}$ & $9.18\\times10^{14}$ \\\\ HD209458b & 0.3 & 1400 & 10 & $7.01\\times10^{19}$ & $4.10\\times10^{17}$ & $2.89\\times10^{15}$ \\\\ HD209458b & 0.3 & 1700 & 1 & $6.97\\times10^{20}$ & $7.67\\times10^{18}$ & $8.02\\times10^{17}$ \\\\ HD209458b & 0.3 & 1700 & 10 & $2.21\\times10^{21}$ & $2.43\\times10^{19}$ & $1.90\\times10^{18}$ \\\\ HD209458b & 0.3 & 2000 & 1 & $1.38\\times10^{22}$ & $3.13\\times10^{20}$ & $4.05\\times10^{19}$ \\\\ HD209458b & 0.3 & 2000 & 10 & $4.52\\times10^{22}$ & $1.05\\times10^{21}$ & $9.42\\times10^{19}$ \\\\ Tres-4b & 0.24 & 2000 & 1 & $6.87\\times10^{22}$ & $2.57\\times10^{21}$ & $1.42\\times10^{20}$ \\\\ Tres-4b & 0.24 & 2250 & 1 & $1.44\\times10^{23}$ & $3.33\\times10^{21}$ & $3.68\\times10^{20}$ \\\\ Tres-4b & 0.24 & 2500 & 1 & $4.62\\times10^{23}$ & $7.87\\times10^{21}$ & $1.54\\times10^{21}$ \\\\ Tres-4b & 0.3 & 2000 & 1 & $4.80\\times10^{22}$ & $9.56\\times10^{20}$ & $3.16\\times10^{19}$ \\\\ Tres-4b & 0.3 & 2250 & 1 & $1.98\\times10^{23}$ & $5.92\\times10^{21}$ & $6.16\\times10^{20}$ \\\\ Tres-4b & 0.3 & 2500 & 1 & $5.13\\times10^{23}$ & $8.75\\times10^{21}$ & $1.55\\times10^{21}$ \\\\ HD189733b & 0.3 & 1500 & 1 & $9.94\\times10^{18}$ & $2.65\\times10^{16}$ & $1.00\\times10^{16}$ \\\\ \\tableline \\end{tabular} \\end{center} \\end{table} We approximate the electrical conductivity profile in the atmosphere with exponential functions: \\begin{eqnarray} \\sigma= \\left\\{\\begin{array}{ll} \\sigma_{\\delta}e^{\\frac{r-(R-\\delta)}{H_{\\delta}}}\\ \\ \\ \\ &\\mbox{$R-\\delta